diff --git "a/batch_s000011.csv" "b/batch_s000011.csv"
new file mode 100644--- /dev/null
+++ "b/batch_s000011.csv"
@@ -0,0 +1,10327 @@
+source,target
+ This case where entropy is a flux function is the natural one to consider when also plasma rotation would be incorporated., This case where entropy is a flux function is the natural one to consider when also plasma rotation would be incorporated.
+ The equation for the density for all three cases can easily be derived by inserting the corresponding pressure equation into the equation for the momentum parallel to the poloidal magnetic field lines given in Eq., The equation for the density for all three cases can easily be derived by inserting the corresponding pressure equation into the equation for the momentum parallel to the poloidal magnetic field lines given in Eq.
+"(9).. The resulting equation is Here, the flux function po corresponds to the density of a related static equilibrium without gravity."," The resulting equation is Here, the flux function $\rho_{0}$ corresponds to the density of a related static equilibrium without gravity."
+" As we will adopt them for the actual stability analysis in our accompanying paper, we briefly discuss the 'straight field line’ coordinates."," As we will adopt them for the actual stability analysis in our accompanying paper, we briefly discuss the `straight field line' coordinates."
+ These coordinates are an essential ingredient of an accurate stability analysis., These coordinates are an essential ingredient of an accurate stability analysis.
+" For the conversion from the Cartesian (x,y,z) to straight field line coordinates (z!=φ,αξJ,a°z), one needs the metric tensor and the Jacobian associated with the non—orthogonal coordinates in which the equilibrium field lines appear to be straight."," For the conversion from the Cartesian $(x,y,z)$ to straight field line coordinates $(x^{1}\equiv\psi, x^{2}\equiv\vartheta, x^{3}\equiv z)$, one needs the metric tensor and the Jacobian associated with the non--orthogonal coordinates in which the equilibrium field lines appear to be straight."
+ Such a transformation is standard practice in MHD stability studies for laboratory tokamak plasmas., Such a transformation is standard practice in MHD stability studies for laboratory tokamak plasmas.
+ The metric elements gj; and the Jacobian J are respectively., The metric elements $g_{ij}$ and the Jacobian $J$ are respectively.
+" Here, the poloidal angle ὁ is constructed such that the magnetic field lines are straight in the (9,z)— plane."," Here, the poloidal angle $\vartheta$ is constructed such that the magnetic field lines are straight in the $(\vartheta, z)$ --plane."
+ The slope of these lines is a flux function where q is the safety factor., The slope of these lines is a flux function where $q$ is the safety factor.
+" Comparing this expression with the one for tokamak plasmas (?),, one should realize that for tokamak plasmas the safety factor q is dimensionless, while here the factor q has a length dimension."," Comparing this expression with the one for tokamak plasmas \citep{Wesson_2004}, one should realize that for tokamak plasmas the safety factor $q$ is dimensionless, while here the factor $q$ has a length dimension."
+" As for tokamak plasmas (???),, we introduce an expression for the poloidal curvature of the magnetic surfaces where the unit vectors n.=Vw/|Vy| and t=Bo/Bo, and By is the poloidal magnetic field."," As for tokamak plasmas \citep{Goedbloed_1975,vanderHolst_2000B,Blokland_2007B}, we introduce an expression for the poloidal curvature of the magnetic surfaces where the unit vectors $\vf{n} = \grad{\psi} / |\grad{\psi}|$ and $\vf{t} = \vf{B}_{\vartheta} / \Btheta$ , and $\Btheta$ is the poloidal magnetic field."
+" The toroidal curvature & that is present in actual tokamak equilibria, is ofcourse zero for a translational symmetric equilibrium."," The toroidal curvature $\kt$ that is present in actual tokamak equilibria, is ofcourse zero for a translational symmetric equilibrium."
+" It is important to realize that the straight field coordinates can only be constructed when the solution v(x,y) has been computed from the extended Grad-Shafranov equation given in Eq."," It is important to realize that the straight field coordinates can only be constructed when the solution $\psi(x,y)$ has been computed from the extended Grad-Shafranov equation given in Eq."
+"(10).. In the previous section, we derived equations for the prominence equilibrium."," In the previous section, we derived equations for the prominence equilibrium."
+" In this section, we quantify the effect of gravity by means of a small gravity expansion."," In this section, we quantify the effect of gravity by means of a small gravity expansion."
+" We will demonstrate in our companion paper that gaps will appear in the continuous MHD spectrum because of mode coupling, which is the result of the presence of gravity."," We will demonstrate in our companion paper that gaps will appear in the continuous MHD spectrum because of mode coupling, which is the result of the presence of gravity."
+" This kind of expansion is similar to thesmall inverse aspect ratio e=a/Ro expansion for tokamak plasmas, where a and Ro are the minor radius of the plasma and the geometry axis of the tokamak, respectively (?).."," This kind of expansion is similar to thesmall inverse aspect ratio $\epsilon = a/R_{0}$ expansion for tokamak plasmas, where $a$ and $R_{0}$ are the minor radius of the plasma and the geometry axis of the tokamak, respectively \citep{Shafranov_1958}."
+" Mathematically, both expansions, small gravity expansion and the small inverse aspect ratio expansion, are Taylor expansions."," Mathematically, both expansions, small gravity expansion and the small inverse aspect ratio expansion, are Taylor expansions."
+" In the remaining part of this paper, we assume the following gravitational potential in which the prominence is embedded where zo is the location of the center of the last closed flux surface of the prominence and the gravity is represented by the constant g."," In the remaining part of this paper, we assume the following gravitational potential in which the prominence is embedded where $x_{0}$ is the location of the center of the last closed flux surface of the prominence and the gravity is represented by the constant $g$."
+ The solarprominence equilibrium is expanded assuming that the gravity is small and that the outer flux surface is circular., The solarprominence equilibrium is expanded assuming that the gravity is small and that the outer flux surface is circular.
+" Using these approximations, the flux surfaces can be represented by slightly displaced circles (?)), which allows for the exploitation of non-orthogonal polar coordinates (r,0,2), where r and 0 are the radius and the polar angle, respectively."," Using these approximations, the flux surfaces can be represented by slightly displaced circles \cite{Shafranov_1958}) ), which allows for the exploitation of non-orthogonal polar coordinates $(r,\theta,z)$, where $r$ and $\theta$ are the radius and the polar angle, respectively."
+" Up to first order, we approximate where A(r) is the Shafranov shift (?)), which is expected to be in the downwards direction and caused by the gravity."," Up to first order, we approximate where $\Delta(r)$ is the Shafranov shift \cite{Shafranov_1958}) ), which is expected to be in the downwards direction and caused by the gravity."
+" As mentioned before, these polar coordinates are non-orthogonal and the associated metric elements are which means that the Jacobian J=r[1—A’ cos(0)]."," As mentioned before, these polar coordinates are non-orthogonal and the associated metric elements are which means that the Jacobian $J \approx r [ 1 - \Delta' \cos(\theta) ]$ ."
+" Using the polar coordinates and expanding the extended Grad-Shafranov equation given in Eqs. (12),, (14),,"," Using the polar coordinates and expanding the extended Grad-Shafranov equation given in Eqs. , ,"
+ and up to first order leads to, and up to first order leads to
+between regions of overlap in separate frames. and the images were averaged.,"between regions of overlap in separate frames, and the images were averaged."
+ At this stage. the separate sets were averaged together to produce the final image.," At this stage, the separate sets were averaged together to produce the final image."
+ The secing was measured from stars in the individual exposures to be in the range L.2aaresec EWIIM. (about 4 pixels) and the registration uncertainties did not cause the ENIM of stars in the final coacdcded images to increase by à measurable amount., The seeing was measured from stars in the individual exposures to be in the range arcsec FWHM (about 4 pixels) and the registration uncertainties did not cause the FWHM of stars in the final coadded images to increase by a measurable amount.
+ Images of photometric standard stars were. taken hroughout the course of each night., Images of photometric standard stars were taken throughout the course of each night.
+ One standard. star was imaged five times in cach filter for each radio galaxy., One standard star was imaged five times in each filter for each radio galaxy.
+ Alultiple coadds were used. both to improve the signal-O-noise ratio. and to ensure that each observation had a similar exposure time to the individual radio galaxy observations.," Multiple coadds were used both to improve the signal-to-noise ratio, and to ensure that each observation had a similar exposure time to the individual radio galaxy observations."
+ In this wav. the point spread function (PSE) of he standard star would sample the same longer-term seeing variations as the galaxy images.," In this way, the point spread function (PSF) of the standard star would sample the same longer-term seeing variations as the galaxy images."
+ Flux calibration solutions were determined: separately for cach night of observation. with the dedispersion for cach night typically being aJ and A and at L' and AL.," Flux calibration solutions were determined separately for each night of observation, with the dispersion for each night typically being at $J$ and $K$ and at $L'$ and $M$."
+ Aperture photometry from our images agrees with that of Lilly Longair (1984) and Lilly. Longair Miller (1985a) to within the quoted uncertainties.," Aperture photometry from our images agrees with that of Lilly Longair (1984) and Lilly, Longair Miller (1985a) to within the quoted uncertainties."
+ In Table 2 we list the photometry measured in a 3-arcsec aperture. which is more appropriate for detecting red nuclear sources.," In Table \ref{tab:phot} we list the photometry measured in a 3-arcsec aperture, which is more appropriate for detecting red nuclear sources."
+ Thermal emission from the telescope is very strong longware of jum. resulting in à vastly increased. background and a large drop in sensitivity.," Thermal emission from the telescope is very strong longward of $\mu$ m, resulting in a vastly increased background and a large drop in sensitivity."
+ Lo addition. the strength. of the quasar nucleus relative to the host galaxy Increases. with wavelength. due to its red. colour.," In addition, the strength of the quasar nucleus relative to the host galaxy increases with wavelength due to its red colour."
+ We therefore. undertake different’ analyses for the short. C/A) and. long (LA) wavelength data., We therefore undertake different analyses for the short ) and long ) wavelength data.
+ In ow J and A images. the host galaxies are well-detected out to. large. radii ancl are likely to. dominate over the nuclei. even in fairly small apertures.," In our $J$ and $K$ images, the host galaxies are well-detected out to large radii and are likely to dominate over the nuclei, even in fairly small apertures."
+ Methods which attempt to estimate the [lux of a nuclear source bv using aperture photometry to correct. small aperture measurcments for starlight are prone to overestimate the strength of the source (Simpson LOOLh) and it is necessary (ο model the galaxy in order to obtain a reliable measurement., Methods which attempt to estimate the flux of a nuclear source by using aperture photometry to correct small aperture measurements for starlight are prone to overestimate the strength of the source (Simpson 1994b) and it is necessary to model the galaxy in order to obtain a reliable measurement.
+ At longer wavelengths. the host galaxy is not detected with &reat significance due to the bright thermal background. and simpler methods can be emploved.," At longer wavelengths, the host galaxy is not detected with great significance due to the bright thermal background, and simpler methods can be employed."
+ To enable the most direct comparison possible of our results with those of T96. we have used an almost identical analysis technique on our data.," To enable the most direct comparison possible of our results with those of T96, we have used an almost identical analysis technique on our data."
+ We refer the reader to that paper fora more detailed description of the method. and for a discussion of its advantages over the more common (and simpler) one-dimensional analysis SSimpson et 11995).," We refer the reader to that paper for a more detailed description of the method, and for a discussion of its advantages over the more common (and simpler) one-dimensional analysis Simpson et 1995)."
+ The two-cimensional approach involves constructing a model image derived from fitting parameters and comparing this with the actual data. using an error frame to allow a quantitative 47 minimization.," The two-dimensional approach involves constructing a model image derived from fitting parameters and comparing this with the actual data, using an error frame to allow a quantitative $\chi^2$ minimization."
+ The assumed model was an elliptical. galaxy obeving a de Vaucouleurs law (de Vaucouleurs 1948) and an unresolved nuclear source., The assumed model was an elliptical galaxy obeying a de Vaucouleurs law (de Vaucouleurs 1948) and an unresolved nuclear source.
+ There are live parameters used to Construct the model: The pixel location of the centre of the galaxy. Gros). was determined using iterative centroiding. and construction of the model galaxy proceeded as follows.," There are five parameters used to construct the model: The pixel location of the centre of the galaxy, $x_{\rm c}$ $y_{\rm
+c}$ ), was determined using iterative centroiding, and construction of the model galaxy proceeded as follows."
+ Each pixel in the image was divided into 400 subpixels. and the flux in each of these subpixels was computed. based on the parameters of the galaxy model.," Each pixel in the image was divided into 400 subpixels, and the flux in each of these subpixels was computed based on the parameters of the galaxy model."
+ This amount of subpixellation ensures that the flux in the central pixel is never uncerestimated, This amount of subpixellation ensures that the flux in the central pixel is never underestimated
+"The local star formation rate volume density pg in a cell is calculated as where py, is the local mass density of molecular hydrogen.",The local star formation rate volume density $\dot{\rho}_\mathrm{S}$ in a cell is calculated as where $\rho_\mathrm{H_2}$ is the local mass density of molecular hydrogen.
+" The star formation time scale is given by where rg=(32Gpa/37)- is the free-fall time, and pa the local gas density (including!/? all hydrogen and helium "," The star formation time scale is given by where $\tau_\mathrm{ff} = (32 G \rho_\mathrm{G} / 3\pi)^{-1/2}$ is the free-fall time, and $\rho_\mathrm{G}$ the local gas density (including all hydrogen and helium species)."
+"The maximum timescale Tmax is set to the free-fall species).time of gas with a hydrogen number density of ng=4-nga2ng,50cm?."," The maximum timescale $\tau_\mathrm{max}$ is set to the free-fall time of gas with a hydrogen number density of $n_\mathrm{H} = n_\mathrm{H\,\textsc{i}}+n_\mathrm{H\,\textsc{ii}}+2 n_\mathrm{H_2} = 50 ~\mathrm{cm}^{-3}$."
+ The star formation efficiency per local free-fall time is set to eg=0.007., The star formation efficiency per local free-fall time is set to $\epsilon_\mathrm{ff} = 0.007$.
+" To ensure that star formation happens only in our numerical analogs of real molecular clouds, we allow star formation only in cells with the molecular mass fraction above fu,=2ng,/ng0.1."," To ensure that star formation happens only in our numerical analogs of real molecular clouds, we allow star formation only in cells with the molecular mass fraction above $f_\mathrm{H_2} = 2 n_\mathrm{H_2}/n_\mathrm{H} = 0.1$."
+" These cells have a range of total gas density from 50 to 104 amu cm""? for the main halo at z=3 (Figure 1)).", These cells have a range of total gas density from 50 to $10^4$ amu $^{-3}$ for the main halo at $z \approx 3$ (Figure \ref{fig:disc}) ).
+ Stellar particles are created via a Poisson process with a characteristic timescale of 2x107 yr., Stellar particles are created via a Poisson process with a characteristic timescale of $2 \times 10^{7}$ yr.
+ This star formation prescription is similar to the recipe SF2 in ?.. , This star formation prescription is similar to the recipe SF2 in \cite{2009ApJ...697...55G}. .
+"Figure shows the disc of the most massive galaxy in our simulation1 at z73, which we call the main halo."," Figure \ref{fig:disc} shows the disc of the most massive galaxy in our simulation at $z \approx 3$, which we call the main halo."
+ The molecular hydrogen forms only in high density regions and hence the stars are confined to these central regions., The molecular hydrogen forms only in high density regions and hence the stars are confined to these central regions.
+" In traditional star formation prescriptions based on the total gas density instead of the molecular hydrogen density, stars would be formed over a much larger volume filled with the lower-density atomic gas."," In traditional star formation prescriptions based on the total gas density instead of the molecular hydrogen density, stars would be formed over a much larger volume filled with the lower-density atomic gas."
+" We ran three versions of the simulation, which are summarized in Table 1.."," We ran three versions of the simulation, which are summarized in Table \ref{tab:simulationsummary}."
+ Simulation A is a full physics run with radiative transfer and non-equilibrium cooling., Simulation A is a full physics run with radiative transfer and non-equilibrium cooling.
+ Simulation Anp is the same as simulation A but without supernova thermal feedback., Simulation $_\mathrm{NF}$ is the same as simulation A but without supernova thermal feedback.
+ Metal enrichment due to supernovae is still included., Metal enrichment due to supernovae is still included.
+ Simulation B is a non-radiative version without cooling and star formation., Simulation B is a non-radiative version without cooling and star formation.
+" In all simulations, the top level |—0 grid is 256? and we allow for up to 9 more refinement levels =9) where each higher level is refined by a factor(Imax 2 with respect to the parent level."," In all simulations, the top level $l=0$ grid is $256^3$ and we allow for up to 9 more refinement levels $l_\mathrm{max}=9$ ) where each higher level is refined by a factor 2 with respect to the parent level."
+ This results in a size of the smallest cells Lo=-279pe , This results in a size of the smallest cells $L_9 = L_\mathrm{box}/(256 \cdot 2^9) = 279 ~\mathrm{pc}$ (comoving).
+"A cell is refined if itsLyox/ dark(256 matter2°) or gas mass(comoving). exceeded 1.07109Mo or 1.3310?Mo, respectively."," A cell is refined if its dark matter or gas mass exceeded $1.07 \times 10^{6} ~\Mo$ or $1.33 \times 10^{5} ~\Mo$, respectively."
+" For the dark matter, this threshold corresponds to the mass of about 6 high resolution particles."," For the dark matter, this threshold corresponds to the mass of about 6 high resolution particles."
+" On each refinement level {, the time step is refined as well according to Av;=Avo/2"", where Avy is the global time step on the top level mesh."," On each refinement level $l$, the time step is refined as well according to $\Delta \nu_l = \Delta \nu_0 / 2^l$, where $\Delta \nu_0$ is the global time step on the top level mesh."
+ The value of Avg is set at the beginning of each top level step so that the Courant-Friedrichs-Lewy condition is fulfilled on all levels (?)..," The value of $\Delta \nu_0$ is set at the beginning of each top level step so that the Courant-Friedrichs-Lewy condition \citep{1928MatAn.100...32C,1967IBMJ...11..215C} is fulfilled on all levels \citep{2002ApJ...571..563K}."
+ In total our simulation (??)A contains 2.89x105 dark matter particles and 3.89x108 gas cells at z=2., In total our simulation A contains $2.89 \times 10^8$ dark matter particles and $3.89 \times 10^8$ gas cells at $z \approx 2$.
+" For the analysis presented in this paper, we mainly concentrate on three snapshots at redshifts around 4, 3, and 2 exact redshifts are 3.76, 2.85, and 2.03) in run A. These (theepochs correspond to 1.69 Gyr, 2.32 Gyr, and 3.29 Gyr after the Big Bang in our cosmology."," For the analysis presented in this paper, we mainly concentrate on three snapshots at redshifts around 4, 3, and 2 (the exact redshifts are 3.76, 2.85, and 2.03) in run A. These epochs correspond to 1.69 Gyr, 2.32 Gyr, and 3.29 Gyr after the Big Bang in our cosmology."
+ T'he output redshifts for different runs match to within Az—0.005., The output redshifts for different runs match to within $\Delta z = 0.005$.
+ Simulation Awp at full resolution was stopped at z=2.77 to save computing time., Simulation $_\mathrm{NF}$ at full resolution was stopped at $z = 2.77$ to save computing time.
+ We ran also a lower-resolution version of Anr with the 8 times more massive dark matter particles but otherwise the same gas physics and parameters., We ran also a lower-resolution version of $_\mathrm{NF}$ with the 8 times more massive dark matter particles but otherwise the same gas physics and parameters.
+ We use the z&2 snapshot from this version in our analysis., We use the $z \approx 2$ snapshot from this version in our analysis.
+" In all snapshots we ran avariant of the Bound Density Maxima halo finder (?) and selected all massive objects with Mago,>10!!Mo in the high resolution region of run A. We then matched simulation A with the other"," In all snapshots we ran avariant of the Bound Density Maxima halo finder \citep{2004ApJ...609..482K} and selected all massive objects with $M_\mathrm{200b} \geq 10^{11}\, \Mo$ in the high resolution region of run A. We then matched simulation A with the other"
+freely available software SExtractor (Bertin&Arnouts 1996)..,freely available software SExtractor \citep{1996A&AS..117..393B}.
+ Although SEsxtractor was developed for the analysis ofoptical data. several authors (Bondietal.2003:Garnct2008a.b:ναetal.2007). have shown that it is able to generate reliable noise maps and locate objects within racio images.," Although SExtractor was developed for the analysis of optical data, several authors \citep{2003A&A...403..857B,2008MNRAS.383...75G,2008MNRAS.387.1037G,2007AJ....133.1331H} have shown that it is able to generate reliable noise maps and locate objects within radio images."
+ It should be noted that SExtractor is highly sensitive to input. parameters and there is no single output. catalogue that will be suitable for all applications., It should be noted that SExtractor is highly sensitive to input parameters and there is no single output catalogue that will be suitable for all applications.
+ We make two radio catalogues., We make two radio catalogues.
+" Firstly we ereate a more conservative ""gold set based on the catalogue of Biges&Lvi- (2006)..", Firstly we create a more conservative `gold' set based on the catalogue of \citet{2006MNRAS.371..963B}.
+" ""Phat work used the 10’.10 VLA pointing alone with an independent reduction. detecting objects at Se in the VLA image alone."," That work used the $10^{\prime}\times10^{\prime}$ VLA pointing alone with an independent reduction, detecting objects at $5\sigma$ in the VLA image alone."
+ The catalogue contains 537 sources. SA of which are contained within the co-acedecd images we use here: the others are either outside the field or resolved away by the much higher. MEBRLLIN resolution.," Their catalogue contains 537 sources, 83 of which are contained within the co-added images we use here; the others are either outside the field or resolved away by the much higher MERLIN resolution."
+ We used SIxtractor in mode to produce an output catalogue of these sources., We used SExtractor in mode to produce an output catalogue of these sources.
+ In order to produce a Larger. fainter catalogue we use Slxtractor in the standard mode to detect islands of (ux above a given threshold.," In order to produce a larger, fainter catalogue we use SExtractor in the standard mode to detect islands of flux above a given threshold."
+ This requires a reliable noise map. which is made by estimating the local background noise at each mesh point of grid across the image (Bertin&Arnouts 1996)..," This requires a reliable noise map, which is made by estimating the local background noise at each mesh point of a grid across the image \citep{1996A&AS..117..393B}."
+ The mesh asize is an important input parameter: if chosen to be too small the background. estimation is allected: by the presence of real objects: if chosen to be too large the small scale variations in the background cannot be reproduced., The mesh size is an important input parameter; if chosen to be too small the background estimation is affected by the presence of real objects; if chosen to be too large the small scale variations in the background cannot be reproduced.
+ We adopt a size of 32 pixels. which corresponds to a2” scale.," We adopt a size of 32 pixels, which corresponds to a scale."
+ We find 3.454]y I rms noise in the resulting map. in close agreement with Muxlowetal. (2005)..," We find $\mu$ Jy $^{-1}$ rms noise in the resulting map, in close agreement with \citet{2005MNRAS.358.1159M}."
+ Using this noise map. sources with a total flux ereater than I05j-]y. (730) were extracted.," Using this noise map, sources with a total flux greater than $\mu$ Jy $(\sim\!\!3\sigma)$ were extracted."
+" This catalogue contains GOL objects and we refer to as the ""silver. set. of objects.", This catalogue contains 691 objects and we refer to as the `silver' set of objects.
+ The optical data we use. forms part of the Great Observatories Origins Deep Survey (GOODS) based on multiband LIST imaging of the WDE and €DE., The optical data we use forms part of the Great Observatories Origins Deep Survey (GOODS) based on multiband HST imaging of the HDF and CDF.
+ The HLDE-N has been imaged in the ACS E435W. FGOGW. FSI4W and FSSOLP bands (D.V./andz respectively) although for the purposes of this work we make use of only the z-band image and catalogue.," The HDF-N has been imaged in the ACS F435W, F606W, F814W and F850LP bands $B, V, i \mbox{ and } z$ respectively) although for the purposes of this work we make use of only the $z$ -band image and catalogue."
+ The 2 band images were observed in 5 epochs separated by 40-50 davs., The $z-$ band images were observed in 5 epochs separated by 40-50 days.
+ In the odd numbered epochs each 10;G ⋅∕⋅fick was tileck: with. a grid. of ⋅⋅3ο individua⋠⋠⋠ ACS pointings., In the odd numbered epochs each $10\arcmin\times16\arcmin$ field was tiled with a grid of $3\times5$ individual ACS pointings.
+ In the even numbered epochs the field was rotated by 45° ancl tiled with 16 separate. pointings., In the even numbered epochs the field was rotated by $45^{\circ}$ and tiled with 16 separate pointings.
+ The band image exposure time was typically 2100s. civicdec into 4 exposures to ensure good cosmic ray rejection., The $z-$ band image exposure time was typically 2100s divided into 4 exposures to ensure good cosmic ray rejection.
+ In each exposure the telescope field of view was shifted by à smal amount to allow optimal sampling of the PSE., In each exposure the telescope field of view was shifted by a small amount to allow optimal sampling of the PSF.
+ The multiple epochs were then combined into a single mosaic., The multiple epochs were then combined into a single mosaic.
+ The observations and image reduction are described in detail in Giavaliseoctal.(2004)..., The observations and image reduction are described in detail in \citet{2004ApJ...600L..93G}.
+ We use the publicly. available 5Extractor configuration files (specified. for each band a /archive.stsci.edu/pub/hlsp/goods/catalog. r2/) το make our catalogues: these configuration [iles have been line tuned to minimise the number of false detections., We use the publicly available SExtractor configuration files (specified for each band at $\mbox{catalog}_{-}$ r2/) to make our catalogues; these configuration files have been fine tuned to minimise the number of false detections.
+ In this section we describe the methods. used. to make estimators of the shear for all of our radio ancl optical SOULCOS., In this section we describe the methods used to make estimators of the shear for all of our radio and optical sources.
+ Lere we summarize the shapelets method which is described more fully in ltefregier(20032)... Relreeier&Bacon(2003) and Massey&Itefregier(2005)...," Here we summarize the shapelets method which is described more fully in \citet{2003MNRAS.338...35R}, , \citet{2003MNRAS.338...48R} and \citet{2005MNRAS.363..197M}. ."
+" In this approach. in the Cartesian formalism. the surface brightness f(x) of a galaxy is decomposed into a series of localisecl orthonormal basis functions D,ss called shapelets: where and where ££5,6]) is the Hermite polvnomial of order m. with the characteristic scale of the basis described. by 3."," In this approach, in the Cartesian formalism, the surface brightness $f(\mathbf{x})$ of a galaxy is decomposed into a series of localised orthonormal basis functions $B_{n_{1},n_{2}}$ called shapelets: where and where $H_{m}(\eta)$ is the Hermite polynomial of order $m$, with the characteristic scale of the basis described by $\beta$ ."
+ Phe series converges most quickly if the characteristic scale 3 ds chosen to be similar to the size of the galaxy. and the centroid of the object is located: accurately.," The series converges most quickly if the characteristic scale $\beta$ is chosen to be similar to the size of the galaxy, and the centroid of the object is located accurately."
+ The sum of ny ancl mo is referred to as the order of the basis functions., The sum of $n_{1}$ and $n_{2}$ is referred to as the order of the basis functions.
+ In practice any decomposition has to be truncated at some order my such that the decomposition vields a sullicienthy accurate model of the galaxy. while also being computationally ellicient. as the computation time of each object's. decompositione Is. AXuias1," In practice any decomposition has to be truncated at some order $n_{\rm max}$ such that the decomposition yields a sufficiently accurate model of the galaxy while also being computationally efficient, as the computation time of each object's decomposition is $\propto n_{\rm max}^4$."
+ From orthonormality. we can find shapelet coellicients for à galaxy by. calculating We use the publicly available shapelets software described in Massev&Relregier(2005) in order to make shapelet decompositionsfor all our objects.," From orthonormality, we can find shapelet coefficients for a galaxy by calculating We use the publicly available shapelets software described in \citet{2005MNRAS.363..197M} in order to make shapelet decompositionsfor all our objects."
+ This code is well tested using optical data (c., This code is well tested using optical data (c.f.
+ Llevmansctal.2006 ancl Masseyetal. 2007)). ancl we seek to extend its applicability to radio data here.," \citealp{2006MNRAS.368.1323H} and \citealp{2007MNRAS.376...13M}) ), and we seek to extend its applicability to radio data here."
+ The code usually fits convolved shapelet cocllicicnts to a galaxy while also optimizing centroid x. 2 and mas using a non-linear algorithm.," The code usually fits convolved shapelet coefficients to a galaxy while also optimizing centroid $\mathbf{x}_{c}$, $\beta$ and $n_{\rm max}$ using a non-linear algorithm."
+ For the radio objects we found that us led to a large number of failures. due to the incorrect stimation of 3 or the centroid wandering olf the edge of 106 postage stamp: in order to stabilise the behaviour. we ix the centroid. position X. to the SExtractor detection 'entroid.," For the radio objects we found that this led to a large number of failures, due to the incorrect estimation of $\beta$ or the centroid wandering off the edge of the postage stamp; in order to stabilise the behaviour, we fix the centroid position $\mathbf{x}_{c}$ to the SExtractor detection centroid."
+ We also have the freedom to fix 3 to 0.4 times 1ο SExtractor ENIM. which we findconsistentlyleads to models with reasonablylow max., We also have the freedom to fix $\beta$ to 0.4 times the SExtractor FWHM which we findconsistentlyleads to models with reasonablylow $n_{\max}$ .
+ Figure | shows some of 1 radio objects and their resulting shapelet models (still convolved with the beam), Figure \ref{fig:shapeletsexamples} shows some of the radio objects and their resulting shapelet models (still convolved with the beam).
+ ‘To deconvolve the beam/PSE from the radio/optical data.," To deconvolve the beam/PSF from the radio/optical data,"
+"SuperWASP-N lightcurve, which comprises 3969 data points obtained over a 118 day period.","SuperWASP-N lightcurve, which comprises 3969 data points obtained over a 118 day period."
+ In the original SuperWASP-N photometry 17 transits were observed with >50% of a transit observed on 10 ocassions., In the original SuperWASP-N photometry 17 transits were observed with $>$ of a transit observed on 10 ocassions.
+" These data led to an ephemeris of 75,—2453139.1748 and P=1.846800 which was used to arrange followup observations.", These data led to an ephemeris of $T_o$ =2453139.1748 and $P$ =1.846800 which was used to arrange followup observations.
+ The transit here has a depth of 0.013 mag., The transit here has a depth of 0.013 mag.
+ and is 137 minutes in duration., and is 137 minutes in duration.
+ WASP-3 was observed with the ccm telescope as part of the Canarian Observatories’ for 2007., WASP-3 was observed with the cm telescope as part of the Canarian Observatories' for 2007.
+ The imaging camera on this telescope has an e2v Technology PLC CCD of 2148x pixels giving a scale of 0.33 arcseconds/pixel and a total field of view of 10.6 arcminutes., The imaging camera on this telescope has an e2v Technology PLC CCD of $2148 \times 2148$ pixels giving a scale of 0.33 arcseconds/pixel and a total field of view of 10.6 arcminutes.
+" Observations were taken during the transit of 2007 August 4, and consist of 327 images of 30and 20 seconds integration in the V and J bands respectively."," Observations were taken during the transit of 2007 August 4, and consist of 327 images of 30and 20 seconds integration in the $V$ and $I$ bands respectively."
+" This night was photometric but suffered from significant Saharan dust extinction, estimated to be ~ 0.4mmag on La Palma from the SuperWASP-N real-time pipeline."," This night was photometric but suffered from significant Saharan dust extinction, estimated to be $\sim 0.4$ mag on La Palma from the SuperWASP-N real-time pipeline."
+ The images were bias subtracted with a stacked bias frame and flat-fielded with a stacked twilight flat field image obtained in both filters using individual flats gathered over the course of the run., The images were bias subtracted with a stacked bias frame and flat-fielded with a stacked twilight flat field image obtained in both filters using individual flats gathered over the course of the run.
+" After the instrumental signatures were removed, source detection and aperture photometry were performed on all science frames using the CASU catalogue extraction software "," After the instrumental signatures were removed, source detection and aperture photometry were performed on all science frames using the CASU catalogue extraction software \citep{il2001}."
+"We chose an aperture size matched to the typical seeing (5 pixels, 1.5"") and selected 5 non-variable comparison stars in the field of WASP-3 to use in deriving the differential photometry."," We chose an aperture size matched to the typical seeing (5 pixels, $^{\prime\prime}$ ) and selected 5 non-variable comparison stars in the field of WASP-3 to use in deriving the differential photometry."
+" For each exposure, we summed the fluxes of the 5 comparison stars and divided by the flux of the target star to derive the differential magnitude of the target."," For each exposure, we summed the fluxes of the 5 comparison stars and divided by the flux of the target star to derive the differential magnitude of the target."
+ The resulting V and I band lightcurves (Figure[I)) of WASP-3 have a precision of ~ 4millimag., The resulting $V$ and $I$ band lightcurves (Figure \ref{fig:lc}) ) of WASP-3 have a precision of $\sim 4$ millimag.
+ Further observations of WASP-3 were made with the Keele University Observatory 60cm Thornton Reflector on 2007 September 10., Further observations of WASP-3 were made with the Keele University Observatory 60cm Thornton Reflector on 2007 September 10.
+" This telescope is equipped with a T65 x 510 pixel Santa Barbara InstrumentGroup (SBIG) ST7 CCD at the f/4.5 Newtonian focus, giving a 0.68 arcsecond/pixel resolution and a 8.63x 5.75 arcminute field of view."," This telescope is equipped with a 765 $\times$ 510 pixel Santa Barbara InstrumentGroup (SBIG) ST7 CCD at the f/4.5 Newtonian focus, giving a 0.68 arcsecond/pixel resolution and a $\times$ 5.75 arcminute field of view."
+ During most of the period the weather was photometric except post egress where some cloud appeared., During most of the period the weather was photometric except post egress where some cloud appeared.
+ Altogether 644x20 sec observations in the R band were obtained., Altogether $\times$ 20 sec observations in the $R$ band were obtained.
+" After applying corrections for bias, dark current and flat fielding in the usual way, aperture photometry on two comparisons were performed using the commercial software AIPAWin "," After applying corrections for bias, dark current and flat fielding in the usual way, aperture photometry on two comparisons were performed using the commercial software AIP4Win \citep{berry2005}."
+Tracking errors and spurious ∙∙electronic noise mean that systematic noise is introduced into the system at an estimated level of 2 millimag with periodicities of 2 and ~20 minutes., Tracking errors and spurious electronic noise mean that systematic noise is introduced into the system at an estimated level of 2 millimag with periodicities of 2 and $\sim$ 20 minutes.
+ No corrections have been applied for this effect., No corrections have been applied for this effect.
+" WASP-3 was observed with the Observatoire de Haute-Provence's mm telescope and the SOPHIE spectrograph (Bouchyetal]2006),, over the 8 nights 2007 July 2 5 and August 27 — 30; a total of 7 usable spectra were acquired."," WASP-3 was observed with the Observatoire de Haute-Provence's m telescope and the SOPHIE spectrograph \citep{b1}, over the 8 nights 2007 July 2 -- 5 and August 27 – 30; a total of 7 usable spectra were acquired."
+ SOPHIE is an environmentally stabilized spectrograph designed to give long-term stability at the level of a few ss!., SOPHIE is an environmentally stabilized spectrograph designed to give long-term stability at the level of a few $^{-1}$.
+" We used the instrument in its high efficiency mode, acquiring simultaneous star and sky spectra through separate fibres with a resolution of R=40000."," We used the instrument in its high efficiency mode, acquiring simultaneous star and sky spectra through separate fibres with a resolution of R=40000."
+" Thorium-Argon calibration images were taken at the start and end of each night, and at 2- to 3-hourly intervals throughout the night."," Thorium-Argon calibration images were taken at the start and end of each night, and at 2- to 3-hourly intervals throughout the night."
+" The radial-velocity drift never exceeded 2-3 m/s, even on a night-to-night basis."," The radial-velocity drift never exceeded 2-3 m/s, even on a night-to-night basis."
+" Conditions during both runs varied from photometric to cloudy, but all nights were affected by strong moonlight."," Conditions during both runs varied from photometric to cloudy, but all nights were affected by strong moonlight."
+" As WASP-3 has magnitude V~10.5, integrations of ssec give a peak signal-to-noise per resolution element of around 40-50."," As WASP-3 has magnitude $V\sim10.5$, integrations of sec give a peak signal-to-noise per resolution element of around 40-50."
+" The 2MASS colours and reduced proper motion for WASP-3 suggest a spectral type of about F7-8V, hence we cross-correlated the spectra against a G2V template provided by the SOPHIE control and reduction software."," The 2MASS colours and reduced proper motion for WASP-3 suggest a spectral type of about F7-8V, hence we cross-correlated the spectra against a G2V template provided by the SOPHIE control and reduction software."
+ In all spectra the cross-correlation functions (CCF) were contaminated by the strong moonlight., In all spectra the cross-correlation functions (CCF) were contaminated by the strong moonlight.
+ We corrected them by using the CCF from the background lights spectrum (mostly the Moon) in the sky fibre., We corrected them by using the CCF from the background light's spectrum (mostly the Moon) in the sky fibre.
+ We then scaled both CCFs using the difference of efficiency between the two fibres., We then scaled both CCFs using the difference of efficiency between the two fibres.
+" Finally we subtracted the corresponding CCF of the background light from the star fibre, and fitted the resulting function by a Gaussian."," Finally we subtracted the corresponding CCF of the background light from the star fibre, and fitted the resulting function by a Gaussian."
+" The parameters obtained allow us to compute the photon-noise uncertainty of the corrected radial velocity measurement (cv), using the relation detailed in (20072): Overall our RV measurements have an average photon-noise uncertainty of 14 m/s. As our radial velocity measurements are not photon-noise limited, we quadratically added a radial velocity component to those uncertainties of about 10 m/s (more details in Section 3.2.1)."," The parameters obtained allow us to compute the photon-noise uncertainty of the corrected radial velocity measurement $\sigma_{RV}$ ), using the relation detailed in \citet{c4}: Overall our RV measurements have an average photon-noise uncertainty of 14 m/s. As our radial velocity measurements are not photon-noise limited, we quadratically added a radial velocity component to those uncertainties of about 10 m/s (more details in Section 3.2.1)."
+ The log of the observations and barycentric RV is given in Table 1., The log of the observations and barycentric RV is given in Table 1.
+" The SOPHIE spectra are individually of modest signal-to-noise, but when summed together they are suitable for a preliminary photospheric analysis of WASP-3."," The SOPHIE spectra are individually of modest signal-to-noise, but when summed together they are suitable for a preliminary photospheric analysis of WASP-3."
+" However, from experience we have found that the SOPHIE standard pipeline reduction does not fully remove the scattered light component within the spectrograph."," However, from experience we have found that the SOPHIE standard pipeline reduction does not fully remove the scattered light component within the spectrograph."
+" While this does not affect radial velocities significantly, it can nonetheless have subtle effects on absorption line depths, adversely affecting the derived spectral synthesis parameters."," While this does not affect radial velocities significantly, it can nonetheless have subtle effects on absorption line depths, adversely affecting the derived spectral synthesis parameters."
+ Therefore we carefully re-reduced the first three raw images taken over 2—5 July 2007 with the echelle data reduction package paying careful attention to the issue of scattered light.," Therefore we carefully re-reduced the first three raw images taken over 2--5 July 2007 with the echelle data reduction package \citep{pv2002}, paying careful attention to the issue of scattered light."
+ ⊓These data are least affected by moonlight., These data are least affected by moonlight.
+" Following our analysis of WASP-1 (Stempels 2007),, we employed the methodology of (2005).., using the same tools, techniques and model atmosphere grid."," Following our analysis of WASP-1 \citep{s2}, , we employed the methodology of \citet{v2}, , using the same tools, techniques and model atmosphere grid."
+ We used the IDL-based software Easy (SME) 1996) to calculate and fit synthetic spectra using a multi-dimensional least squares approach., We used the -based software ) \citep{v1} to calculate and fit synthetic spectra using a multi-dimensional least squares approach.
+sriehtest Cluster Galaxies (BCCis) include the most massive ealaxies in the Universe.,Brightest Cluster Galaxies (BCGs) include the most massive galaxies in the Universe.
+ Models. of hierarchical structure formation naturally feature the ongoing growth of the most massive galaxies bv mergers (Peebles&Yu1970).. and BCCs are predicted to have undergone more mergers than loss massive galaxies (e.g. DeLucia&Dlaizot2007..," Models of hierarchical structure formation naturally feature the ongoing growth of the most massive galaxies by mergers \citep{peebles70}, and BCGs are predicted to have undergone more mergers than less massive galaxies (e.g. \citealt{delucia07}."
+ Thus determining the merging history of BCCs is a particularly sensitive test of current formation models., Thus determining the merging history of BCGs is a particularly sensitive test of current formation models.
+ As the analvtic models of Bowerctal.(2006):Crotonet(2006) show. massive galaxies are over-produced in. N-body. dark matter cosmological simulations. and [feedback mechanisms are required to bring the luminosity function into agreement with observations.," As the semi-analytic models of \cite{bower06,croton06} show, massive galaxies are over-produced in N-body dark matter cosmological simulations, and feedback mechanisms are required to bring the luminosity function into agreement with observations."
+ Tracing the recent assembly history of BC's is vital to future development of galaxy formation models., Tracing the recent assembly history of BCGs is vital to future development of galaxy formation models.
+ Observationally the evidence for X€ erowth is contraclictorv: Studies of the luminosities and stellar masses of BCCGs show little evolution in mass since z~115 (e.g. Broughetal.2002:Collins 2009)) and their steep metallicity gracicnts are consistent with passive evolution since 2— (Broughetal.2007).," Observationally the evidence for BCG growth is contradictory: Studies of the luminosities and stellar masses of BCGs show little evolution in mass since $z\sim1-1.5$ (e.g. \citealt{brough02, collins09}) ) and their steep metallicity gradients are consistent with passive evolution since $z\sim2$ \citep{brough07}."
+. However. the large radii ancl low surface brightnesses of BCCs compared to normal elliptical galaxies are consistent with products. of major. dissipationless mergers (c.g. Ocgerle&Hoessel1991:Broughetal.2005:vonderLindenct2007:Lauer 2007)).," However, the large radii and low surface brightnesses of BCGs compared to normal elliptical galaxies are consistent with products of major, dissipationless mergers (e.g. \citealt{oegerle91,brough05, vDL07, lauer07}) )."
+" Their sizes ancl velocity. dispersions may have also evolved faster than less-massive early-type galaxies since z0.3 (Bernardi2009 although. οἱ, Stottetal. 2011))."," Their sizes and velocity dispersions may have also evolved faster than less-massive early-type galaxies since $z\sim0.3$ \citealt{bernardi09} although, c.f. \citealt{stott11}) )."
+ While BCGs are frequently observed to have multiple nuclei and close companions (c.g. Schneideretal. 1983)). there are only," While BCGs are frequently observed to have multiple nuclei and close companions (e.g. \citealt{schneider83}) ), there are only"
+the accretion disk emission when the svnchrotron emission decreased (Abdoetal.2009b).,the accretion disk emission when the synchrotron emission decreased \citep{abd09b}.
+. It is likely also the case in X-ray waveband (e.g.Foschinietal.2009a).. (hough it is elaimed that in most radio loud quasars the contribution from a hot disk corona to the observed N-ravs is jeglieible (in the hard. X-ray band >2 keV). except for the steep-spectrum soft. X-ray excess below 1 keV (e.g.. Brinkmann et al.," It is likely also the case in X-ray waveband \cite[e.g.][]{fos09a}, though it is claimed that in most radio loud quasars the contribution from a hot disk corona to the observed X-rays is negligible (in the hard X-ray band $>$ 2 keV), except for the steep-spectrum soft X-ray excess below 1 keV (e.g., Brinkmann et al."
+ 1997; Yuan el al., 1997; Yuan et al.
+ 2000)., 2000).
+ Indeed. the integrated model including svnchrotron jet emission. disk-corona emission and inverse Compton emission [rom jet has been used to model SED of NLSIs. trom which the nature of radio loud NLS1s can be well studied (Abdoetal.2009a.b.c:Fosehini2009b).," Indeed, the integrated model including synchrotron jet emission, disk-corona emission and inverse Compton emission from jet has been used to model SED of NLS1s, from which the nature of radio loud NLS1s can be well studied \citep{abd09a,abd09b,abd09c,fos09b}."
+. While it strongly confirms the presence of a relativistic jet in radio-loud NLSIs. the —rav detection is also important to study the jet properties bx modeling the SEDs. e.g. the jet power. from which the characteristic of radio-loud NLS1s can be explored.," While it strongly confirms the presence of a relativistic jet in radio-loud NLS1s, the $\gamma-$ ray detection is also important to study the jet properties by modeling the SEDs, e.g. the jet power, from which the characteristic of radio-loud NLS1s can be explored."
+ Through the model fit including svuchrotron sell-Compton (5C). external Compton (EC). and accretion disk-corona in four gamuna-ray detected NLSIs. their jet powers are found in (he average range of blazars with (wo sources in (he region of quasars. and another (wo in the range of ivpical of BL Lac objects (Abdoοἱal.2009€).," Through the model fit including synchrotron self-Compton (SSC), external Compton (EC), and accretion disk-corona in four gamma-ray detected NLS1s, their jet powers are found in the average range of blazars with two sources in the region of quasars, and another two in the range of typical of BL Lac objects \citep{abd09c}."
+. ILowever. the main differences with respect to blazars are in the black hole masses and accretion rates. as argued in Abdoetal.(2009€). with the former about. 1-2 orders of magnitude lower than the (vpical blazar masses. aud the later obviously higher than (hose of blazars.," However, the main differences with respect to blazars are in the black hole masses and accretion rates, as argued in \cite{abd09c}, with the former about 1-2 orders of magnitude lower than the typical blazar masses, and the later obviously higher than those of blazars."
+ Moreover. blazars are usually hosted by elliplical galaxies. while it is likely to be spiral ones in raclio-loud NLSIs (e.g.Zhouetal. 2006).," Moreover, blazars are usually hosted by elliptical galaxies, while it is likely to be spiral ones in radio-loud NLS1s \cite[e.g.][]{zho06}."
+. From these observational eviclences. Abdoetal.(20090). claimed. that racio-Ioud NLS1s may represent a third subset of 5—rav AGNs. besides blazars and radio galaxies.," From these observational evidences, \cite{abd09c} claimed that radio-loud NLS1s may represent a third subset of $\gamma-$ ray AGNs, besides blazars and radio galaxies."
+ If this is (hie case. il remains unclear whether radio-oud NLS1s should follow the blazar sequence (e.g. RAS J1629024-4007 in Chis work).," If this is the case, it remains unclear whether radio-loud NLS1s should follow the blazar sequence (e.g. RXS J16290+4007 in this work)."
+ This certainly needs Iurther investigations., This certainly needs further investigations.
+ On the other hand. it is still not clear why racio-loud NLS1s host a relativistic jet. aud Low ib is formed.," On the other hand, it is still not clear why radio-loud NLS1s host a relativistic jet, and how it is formed."
+ It can be even more complicated in terms of the fact that the host galaxies ol NLSIs is generally of spiral (wpe. which breaks the paradigm: associating relativistic jets with giant elliptical (e.g.Marscher2009).," It can be even more complicated in terms of the fact that the host galaxies of NLS1s is generally of spiral type, which breaks the paradigm associating relativistic jets with giant elliptical \cite[e.g.][]{mars09}."
+. Although the accretion disk and jet are found to be closely related (e.g. 2009a).. the details of disk-jet coupling is not known vet. besides that jet formation and radio loud/cquiet clichotonw of AGNs are not well understood 2009).," Although the accretion disk and jet are found to be closely related \cite[e.g.][]{cao01,gu09a}, the details of disk-jet coupling is not known yet, besides that jet formation and radio loud/quiet dichotomy of AGNs are not well understood \cite[e.g.][]{tch09}."
+. As one possibility. jel activity can be intermittent. due to. lor example. the accretion disk instability (e.g.Czernyetal.2009).. which is recently adopted to explain CSS/GPS sources (Wu2009a).," As one possibility, jet activity can be intermittent, due to, for example, the accretion disk instability \cite[e.g.][]{cze09}, which is recently adopted to explain CSS/GPS sources \citep{wu09a}."
+. The existence of double double radio sources seems {ο support the intermittent scenario (e.g.Marecki&Szablewski2009)., The existence of double double radio sources seems to support the intermittent scenario \cite[e.g.][]{mar09}.
+. Optically. NLSIs are thought to be voung AGNs with small black hole mass accreting at high accretion rate. implving the central accretion process are al the early stage of accretion history.," Optically, NLS1s are thought to be young AGNs with small black hole mass accreting at high accretion rate, implying the central accretion process are at the early stage of accretion history."
+ In. radio band. the compact nature of radio structure of CSS sources are believed to be due to the [act," In radio band, the compact nature of radio structure of CSS sources are believed to be due to the fact"
+"is approximately linear, the slope is shallower than expected.","is approximately linear, the slope is shallower than expected."
+" An accurate calibration of o is important, because in turn it is used as an input for the evaluation and fitting of the warp diffusion coefficient a2."," An accurate calibration of $\alpha$ is important, because in turn it is used as an input for the evaluation and fitting of the warp diffusion coefficient $\alpha_2$."
+" The disagreement found in ? prompted us to examine the method used to calibrate o in greater detail, resulting in our implementation of the fitting procedure described in Section 4.2 — essentially a quantitative version of the procedure performed in ?.."," The disagreement found in \citetalias{LP07} prompted us to examine the method used to calibrate $\alpha$ in greater detail, resulting in our implementation of the fitting procedure described in Section \ref{sec:fit} — essentially a quantitative version of the procedure performed in \citetalias{LP07}."
+" In considering this issue, we have also explored the effect of the inner boundary condition of the 1D disc evolution on the measurement of a."," In considering this issue, we have also explored the effect of the inner boundary condition of the 1D disc evolution on the measurement of $\alpha$."
+" Indeed, the main feature which is used for the"," Indeed, the main feature which is used for the"
+"emission), are variable at different wavelengths.","emission), are variable at different wavelengths."
+ Only and of the sources above the black-body line and below the power-law line. respectively. show variability.," Only and of the sources above the black-body line and below the power-law line, respectively, show variability."
+ One aim of this paper is to study the possible evolutionary stages of the AGN in the CJF sample., One aim of this paper is to study the possible evolutionary stages of the AGN in the CJF sample.
+ Different subsamples of the CJF. and the individual sources within them. can probe these different evolutionary stages.," Different subsamples of the CJF, and the individual sources within them, can probe these different evolutionary stages."
+ Mergers. starburst activity. and BBH systems are integral parts of this study.," Mergers, starburst activity, and BBH systems are integral parts of this study."
+ One of the most prominent examples supporting the link between starburst and BBH systems 1s NGC 6240. one of few sources directly observed to have a binary core (?)).," One of the most prominent examples supporting the link between starburst and BBH systems is NGC 6240, one of few sources directly observed to have a binary core \citealt{Komossa2003}) )."
+ This system ts an ULIRG. hosts two AGN. and also clearly exhibits ongoing starburst activity. making it an archetype of the evolutionary scenario discussed here.," This system is an ULIRG, hosts two AGN, and also clearly exhibits ongoing starburst activity, making it an archetype of the evolutionary scenario discussed here."
+ By selecting sources with a companion. one can investigate the earliest stages of the merging process.," By selecting sources with a companion, one can investigate the earliest stages of the merging process."
+ Accordingly. by selecting the sources with disturbed morphologies and large infrared fluxes. we probe an intermediate phase (e.g.. Mrk 231).," Accordingly, by selecting the sources with disturbed morphologies and large infrared fluxes, we probe an intermediate phase (e.g., Mrk 231)."
+ Finally. sources with relaxed morphologies but almost periodic variability at multiple wavelengths are likely to represent systems in. which an assumed BBH has sunk to the center of the system.," Finally, sources with relaxed morphologies but almost periodic variability at multiple wavelengths are likely to represent systems in which an assumed BBH has sunk to the center of the system."
+ ? conducted simulations of equal-mass gas-rich mergers. classifying the sources into six distinct merging phases: pre-merger. first pass. maximal separation. merger. post-merger. and remnant.," \citet{Lotz2008} conducted simulations of equal-mass gas-rich mergers, classifying the sources into six distinct merging phases: pre-merger, first pass, maximal separation, merger, post-merger, and remnant."
+ Using this classification scheme for our sources. we classify objects into pre and post-merger stages.," Using this classification scheme for our sources, we classify objects into pre and post-merger stages."
+ The incompleteness of the information. available. about our sample does not allow us to make unambiguous claims for the classification of some sources (especially when differentiating between for example either. pre-merger and maximal separation. or post-merger and remnant classes).," The incompleteness of the information available about our sample does not allow us to make unambiguous claims for the classification of some sources (especially when differentiating between for example either pre-merger and maximal separation, or post-merger and remnant classes)."
+ Among the distorted sources. those with companions can be classified as either a pre-merger or maximal separation phase. while those with no companions can be categorized as being in either the first pass. the post-merger. or remnant phase.," Among the distorted sources, those with companions can be classified as either a pre-merger or maximal separation phase, while those with no companions can be categorized as being in either the first pass, the post-merger, or remnant phase."
+ For the last group. this translates. roughly. to an age of 2- Gyr after the initial approach of the progenitor systems.," For the last group, this translates, roughly, to an age of 2-4 Gyr after the initial approach of the progenitor systems."
+ Exceptions to the above constitute four sources (3C. 84. Mrk 231. 3C 371. and Mrk 501) that exhibit both distorted morphologies and considerable starburst activity (e.g. 2:; 2)).," Exceptions to the above constitute four sources (3C 84, Mrk 231, 3C 371, and Mrk 501) that exhibit both distorted morphologies and considerable starburst activity (e.g., \citealt{Richards2005}; \citealt{Condon2002}) )."
+ These characteristics put them in the merger phase and would constrain their age to be «2 Gyr after the initial approach., These characteristics put them in the merger phase and would constrain their age to be $<2$ Gyr after the initial approach.
+ In Table 10 we indicate the 27 most suitable representatives of different evolutionary phases (following ?))., In Table \ref{tab:evolution_candidates} we indicate the 27 most suitable representatives of different evolutionary phases (following \citealt{Lotz2008}) ).
+ Sources with companions but no sign of disturbed morphologies. are selected as pre-merger systems.," Sources with companions but no sign of disturbed morphologies, are selected as pre-merger systems."
+ Sources with disturbed morphologies but no companions and no apparent starburst activity. are assigned to the first pass phase when the two systems are at a minimum separation for the first time.," Sources with disturbed morphologies but no companions and no apparent starburst activity, are assigned to the first pass phase when the two systems are at a minimum separation for the first time."
+ The next phase corresponds to that of the maximal separation between the merging systems., The next phase corresponds to that of the maximal separation between the merging systems.
+ For this phase. sources with disturbed morphologies and detected companions are selected.," For this phase, sources with disturbed morphologies and detected companions are selected."
+ Sources identified as undergoing the first pass may also be merger remnants., Sources identified as undergoing the first pass may also be merger remnants.
+ More information is required to make this distinction., More information is required to make this distinction.
+MCG-6-30-15 (Miniutti et al.,MCG–6-30-15 (Miniutti et al.
+ 2007)., 2007).
+" For this reason, it is not possible to constrain simultaneously all the line parameters (see Table 1)): these include the disc inner radius (rin) and inclination with respect to the line of sight and the emissivity index q (under the assumption of a radial emissivity profile e(r)οςr~%)."," For this reason, it is not possible to constrain simultaneously all the line parameters (see Table \ref{t1}) ): these include the disc inner radius $r_\rmn{in}$ ) and inclination with respect to the line of sight and the emissivity index $q$ (under the assumption of a radial emissivity profile $\epsilon (r) \propto r^{-q}$ )."
+" None the less, we note that the best-fitting inner radius of ~13 rg does not necessarily require the extreme gravity regime typical of rapidly rotating black holes."," None the less, we note that the best-fitting inner radius of $\sim$ 13 $r_\rmn{g}$ does not necessarily require the extreme gravity regime typical of rapidly rotating black holes."
+" It is important to stress that alternative explanations invoking a spectral upturn due to complex absorption effects, such as those proposed for the same MCG-6-30-15, are definitely not viable in the case of Ark 120."," It is important to stress that alternative explanations invoking a spectral upturn due to complex absorption effects, such as those proposed for the same MCG–6-30-15, are definitely not viable in the case of Ark 120."
+" 4) In the latter stage, the properties of the two narrow lines are exactly the same as found in the original model."," 4) In the latter stage, the properties of the two narrow lines are exactly the same as found in the original model."
+" Thus, it is worth investigating another possibility by retaining the basic template and allowing the width of the two lines to vary."," Thus, it is worth investigating another possibility by retaining the basic template and allowing the width of the two lines to vary."
+" The consequent fit refinement is not as significant as when a disc line is involved, being now x2/d.o.f.=0.735/282: the difference of Ax?~—15.2 obtained with the loss of two degrees of freedom is not likely to be simply a chance improvement and supports the relativistic line detection."," The consequent fit refinement is not as significant as when a disc line is involved, being now $\chi^2_\nu/\rmn{d.o.f.}=0.735/282$: the difference of $\Delta \chi^2 \simeq -15.2$ obtained with the loss of two degrees of freedom is not likely to be simply a chance improvement and supports the relativistic line detection."
+" Moreover, although the double-peaked energy of the blended feature is still consistent within the errors with neutral and H-like iron emission (Table 1)), the resulting width of c=113(+21) eV poses the question about the physical location wherein these lines arise."," Moreover, although the double-peaked energy of the blended feature is still consistent within the errors with neutral and H-like iron emission (Table \ref{t1}) ), the resulting width of $\sigma = 113 (\pm 21)$ eV poses the question about the physical location wherein these lines arise."
+" Such a value, in fact, corresponds to a full width at half-maximum (FWHM) broadening of ~12x10? km s!, which is a factor of —2 larger than that observed in the optical permitted lines (FWHM Hf c5800 km s7!; Wandel, Peterson Malkan 1999)."," Such a value, in fact, corresponds to a full width at half-maximum (FWHM) broadening of $\sim 12 \times 10^3$ km $^{-1}$, which is a factor of $\sim$ 2 larger than that observed in the optical permitted lines (FWHM $\beta$ $\simeq 5800$ km $^{-1}$; Wandel, Peterson Malkan 1999)."
+ This could hint at a sort of X-ray broad-line region (BLR) internal to the optical one., This could hint at a sort of X-ray broad-line region (BLR) internal to the optical one.
+" Any further discussion on the possible origin of these lines (either broad or narrow) is deferred to the next section, in which we address this issue within the context of X-ray reflection models."," Any further discussion on the possible origin of these lines (either broad or narrow) is deferred to the next section, in which we address this issue within the context of X-ray reflection models."
+" We are confident that relativistic effects are in place and have to be taken into account also when considering the entire spectral range, since all the interpretations of the iron emission profile in Ark 120 making no resort to a disc line turn out to be less successful and convincing."," We are confident that relativistic effects are in place and have to be taken into account also when considering the entire spectral range, since all the interpretations of the iron emission profile in Ark 120 making no resort to a disc line turn out to be less successful and convincing."
+" We now extend our analysis to the whole 0.640 keV energy range, in order to understand the origin of the spectral Excess emission beyond ~20-30 keV is usually interpreted as due to the reprocessing of the primary X-ray radiation: the combination of photoelectric absorption and Compton scattering in the illuminated material gives rise to a broad reflection hump (e.g. George Fabian 1991, and references therein)."," We now extend our analysis to the whole 0.5–40 keV energy range, in order to understand the origin of the spectral Excess emission beyond $\sim$ 20–30 keV is usually interpreted as due to the reprocessing of the primary X-ray radiation: the combination of photoelectric absorption and Compton scattering in the illuminated material gives rise to a broad reflection hump (e.g. George Fabian 1991, and references therein)."
+" Besides this additional continuum component and iron fluorescence, below ~2 keV the reflected spectrum is expected to be dominated by a wealth of emission lines from oxygen and other abundant elements, like C, N, Ne, Mg, Si, S (Ross Fabian 1993)."," Besides this additional continuum component and iron fluorescence, below $\sim$ 2 keV the reflected spectrum is expected to be dominated by a wealth of emission lines from oxygen and other abundant elements, like C, N, Ne, Mg, Si, S (Ross Fabian 1993)."
+" Since it is fairly conceivable that in many cases the accretion flow itself acts as the most efficientmirror, depending on the ionization stage of the disc outer layers and on the relativistic motions of the inner regions, the stack of individual features can be blurred into the smooth shape of the soft excess."," Since it is fairly conceivable that in many cases the accretion flow itself acts as the most efficient, depending on the ionization stage of the disc outer layers and on the relativistic motions of the inner regions, the stack of individual features can be blurred into the smooth shape of the soft excess."
+" As stated above, the slight spectral curvature observed at ~6 keV in Ark 120 cannot be explained as the product of (multiple) covering effects, suggesting instead the presence of a broad skewed profile within the iron emission feature and indicating a strong gravity regime."," As stated above, the slight spectral curvature observed at $\sim$ 6 keV in Ark 120 cannot be explained as the product of (multiple) covering effects, suggesting instead the presence of a broad skewed profile within the iron emission feature and indicating a strong gravity regime."
+" It is therefore reasonable to include in our general model two reflection components, for which we have used the self-consistent table models of Ross Fabian (2005): the first one is expected to arise from almost neutral material at great distance from the X-ray emitting region, likely situated on the dusty torus scale."," It is therefore reasonable to include in our general model two reflection components, for which we have used the self-consistent table models of Ross Fabian (2005): the first one is expected to arise from almost neutral material at great distance from the X-ray emitting region, likely situated on the dusty torus scale."
+" The second one can be ascribed to the partially ionized surface of the disc, and has been convolved with the kernel in to account for relativistic effects."," The second one can be ascribed to the partially ionized surface of the disc, and has been convolved with the kernel in to account for relativistic effects."
+" The model obviously comprises the primary power law, absorbed by the Galactic column density, and also the two unresolved gaussian lines identified above, corresponding to neutral and H-like iron emission."," The model obviously comprises the primary power law, absorbed by the Galactic column density, and also the two unresolved gaussian lines identified above, corresponding to neutral and H-like iron emission."
+" Galactic absorption is found to be consistent with Nu=0.98x107?! cm?, and has been frozen accordingly."," Galactic absorption is found to be consistent with $N_\rmn{H}=0.98 \times 10^{-21}$ $^{-2}$, and has been frozen accordingly."
+" The reflection templates include among their free parameters iron abundance (which has been assumed to be single-valued across the source), ionization state (described by the quantity £=47F/ny, where ng is the hydrogen number density of the gas exposed to an X-ray flux F) and photon index of the illuminating power law (which has been taken to be one and the same with the direct power law)."," The reflection templates include among their free parameters iron abundance (which has been assumed to be single-valued across the source), ionization state (described by the quantity $\xi = 4 \pi F/n_\rmn{H}$, where $n_\rmn{H}$ is the hydrogen number density of the gas exposed to an X-ray flux $F$ ) and photon index of the illuminating power law (which has been taken to be one and the same with the direct power law)."
+particle’ for each cluster that initially moves with the average velocity of the gas.,particle' for each cluster that initially moves with the average velocity of the gas.
+" On scales =100 pc this effect plays only a small role presumably, as the typical distance that stars travel within 20 Myr is of the order of ~100 pc (assuming an rms velocity of ~5 km s~')."," On scales $\gtrsim{}100$ pc this effect plays only a small role presumably, as the typical distance that stars travel within 20 Myr is of the order of $\sim{}100$ pc (assuming an rms velocity of $\sim{}5$ km $^{-1}$ )."
+ Some scatter on large scales may arise from high-velocity run-away stars (Blaauw1961;Stone ," Some scatter on large scales may arise from high-velocity run-away stars \citep{1961BAN....15..265B, 1991AJ....102..333S}."
+The decline of the scatter 1991)..with increasing averaging scale is obviously related to the spatial averaging over a larger number of resolution elements Nyes., The decline of the scatter with increasing averaging scale is obviously related to the spatial averaging over a larger number of resolution elements $N_{\rm res}$.
+" Naively, we would expect a scaling proportional to 1/9Νιος, where Nus«d? if the Ha is filling the volume relatively uniformly, or ος1? if the H» is confined to a disk."," Naively, we would expect a scaling proportional to $1/\sqrt{N_{\rm res}}$, where $N_{\rm res}\propto{}l^3$ if the $\H2$ is filling the volume relatively uniformly, or $\propto{}l^2$ if the $\H2$ is confined to a disk."
+" However, the scale dependence shown in Fig."," However, the scale dependence shown in Fig."
+ 4 seems to be much shallower., \ref{fig:scaleScatter} seems to be much shallower.
+" In order to understand this better, we now include a log-normal (intrinsic) scatter of the in equation with ranging from 0.1 to 1 dex on the 65 pe (1)),scale."," In order to understand this better, we now include a log-normal (intrinsic) scatter of the in equation \ref{eq:SFR}) ), with $\sigma/\ln{}10$ ranging from 0.1 to 1 dex on the 65 pc scale."
+ Fig., Fig.
+ 5 c/1n10shows that this scatter decreases with increasing averaging scale roughly as a power law l? with exponent α=0.5., \ref{fig:scaleScatter2} shows that this scatter decreases with increasing averaging scale roughly as a power law $\sigma_l\propto{}l^{-\alpha}$ with exponent $\alpha\approx{}0.5$.
+" As we discuss in more detail in Appendix B the rather gradual decline is caused by both the finite width c, of the Hy density distribution and the geometrical arrangement of the molecular hydrogen.", As we discuss in more detail in Appendix \ref{sect:dep} the rather gradual decline is caused by both the finite width $\sigma_\rho$ of the $\H2$ density distribution and the geometrical arrangement of the molecular hydrogen.
+" In particular, we find that a sensible value σρ©1.5 (McKee&Ostriker2007) and a two-dimensional (2D, disk) arrangement of the molecular hydrogen naturally leads to ae0.5."," In particular, we find that a sensible value $\sigma_\rho\approx{}1.5$ \citep{2007ARA&A..45..565M} and a two-dimensional (2D, disk) arrangement of the molecular hydrogen naturally leads to $\alpha\approx{}0.5$."
+" On the other hand, a purely one-dimensional (1D) arrangement of the Hg is only consistent with az0.5 if the width in the density distribution is very small («&1), while a three-dimensional (3D) configuration would require σρ72.5."," On the other hand, a purely one-dimensional (1D) arrangement of the $\H2$ is only consistent with $\alpha\approx{}0.5$ if the width in the density distribution is very small $\ll{}1$ ), while a three-dimensional (3D) configuration would require $\sigma_\rho\approx{}2.5$."
+" We have also tested that the exponent decreases if the inserted intrinsic scatter becomes large, ie. o>1."," We have also tested that the exponent decreases if the inserted intrinsic scatter becomes large, i.e., $\sigma\gg{}1$."
+" In fact, the top curves (c=2.3) in Fig."," In fact, the top curves $\sigma=2.3$ ) in Fig."
+ 5 clearly show this flattening., \ref{fig:scaleScatter2} clearly show this flattening.
+" Since the Πο density distribution depends on the metallicity and radiation field (see Fig. 2)),"," Since the $\H2$ density distribution depends on the metallicity and radiation field (see Fig. \ref{fig:regfit}) ),"
+" α should depend on it, too."," $\alpha$ should depend on it, too."
+" With the data points shown in Fig. 5,,"," With the data points shown in Fig. \ref{fig:scaleScatter2},"
+" we obtain a=0.52+0.04 for Z=0.1, Umw=100, and a=0.43€0.04 for Z=1, Umw=0.1."," we obtain $\alpha=0.52\pm{}0.04$ for $Z=0.1$, $\UMW=100$, and $\alpha=0.43\pm{}0.04$ for $Z=1$, $\UMW=0.1$ ."
+ Let us now consider the case in which a scatter σι is inserted on scale | (we assume a set of discrete scales that change by a factor 2)., Let us now consider the case in which a scatter $\tilde{\sigma}_l$ is inserted on scale $l$ (we assume a set of discrete scales that change by a factor 2).
+" If the different scatter contributions add in quadrature, the total scatter σι on a given scale { is simply given by With knowledge of a, this equation allows the computation of the amount of scatter σι that is introduced on scale { from the measurement of the scatter on scales { and 1/2."," If the different scatter contributions add in quadrature, the total scatter $\sigma_l$ on a given scale $l$ is simply given by With knowledge of $\alpha$, this equation allows the computation of the amount of scatter $\tilde{\sigma}_l$ that is introduced on scale $l$ from the measurement of the scatter on scales $l$ and $l/2$."
+" Presumably, different physical mechanisms may introduce different amounts of scatter on different scales."," Presumably, different physical mechanisms may introduce different amounts of scatter on different scales."
+ Studying the scale dependence of the scatter may therefore be helpful to uncover the responsible physical mechanism(s)., Studying the scale dependence of the scatter may therefore be helpful to uncover the responsible physical mechanism(s).
+" The scatter in the relation on the scale of ~100 pc has been attributed to the evolution of molecular clouds over their life time (see, e.g., Onoderaetal.2010;Schruba 2010))."," The scatter in the relation on the scale of $\sim{}100$ pc has been attributed to the evolution of molecular clouds over their life time (see, e.g., \citealt{2010arXiv1009.1971O, 2010arXiv1009.1651S}) )."
+" In this picture, young molecular clouds have not yet formed stars, but contain large amounts of Πο and hence fall “below” the average relation."," In this picture, young molecular clouds have not yet formed stars, but contain large amounts of $\H2$ and hence fall “below” the average relation."
+" On the other hand, clouds that are near the end of their lives are heavily star forming and/or have lost some fraction of their molecular hydrogen, hence they lie “above” the relation."," On the other hand, clouds that are near the end of their lives are heavily star forming and/or have lost some fraction of their molecular hydrogen, hence they lie “above” the relation."
+ This picture cannot be reconciled with an Ho-based star formation law of the form of equation as long as the gas consumption time scale TsrR/esrn is (1))treated as a constant., This picture cannot be reconciled with an $\H2$ -based star formation law of the form of equation \ref{eq:SFR}) ) as long as the gas consumption time scale $\tau_{\rm SFR}/\epsilon_{\rm SFR}$ is treated as a constant.
+" Hence, this explanation of the scatter in the relation implies that has to be a time-dependent quantity."," Hence, this explanation of the scatter in the relation implies that $\epsilon_{\rm SFR}/\tau_{\rm SFR}$ has to be a time-dependent quantity."
+" If Tspg is espn/Tsrmapproximately constant, then the star formation efficiency will need to change over the life time of a molecular cloud (e.g., Murray2011,, but see Feldmann&Gnedi"," If $\tau_{\rm SFR}$ is approximately constant, then the star formation efficiency will need to change over the life time of a molecular cloud (e.g., \citealt{2010arXiv1007.3270M}, but see \citealt{FeldmannM}) )."
+n Our interpretation is 2011)).different., Our interpretation is different.
+" We show that a large amount of scatter in the relation can be explained by the fact that observations do not measure the instantaneous rate of star formation, butrather the number of stars that formed within a finite time interval in the past."," We show that a large amount of scatter in the relation can be explained by the fact that observations do not measure the instantaneous rate of star formation, butrather the number of stars that formed within a finite time interval in the past."
+ Our numerical models predict that the scatter seen on scales of ~100 pc should be small, Our numerical models predict that the scatter seen on scales of $\sim{}100$ pc should be small
+"Integrating dT from the shock towards the upstream, the value of Γ/Τ,, (for values larger than 1/T,) has a universal profile (independent of T,,) as a function of n,ozz'(14-2]n|T2])/T2.","Integrating $d\G$ from the shock towards the upstream, the value of $\G/\Gu$ (for values larger than $1/\Gu$ ) has a universal profile (independent of $\G_u$ ) as a function of $n_u \sigma_T
+z'(1+2\ln[\Gu^2])/\Gu^2$."
+ 'This profile is presented in figure 3.., This profile is presented in figure \ref{fig3}.
+" This profile implies that the shock width in unit of pair unloaded Thomson optical depth in the shock frame isroughly!!,, οςI?."," This profile implies that the shock width in unit of pair unloaded Thomson optical depth in the shock frame is, $\propto \G_u^2$."
+" The value of the proportionality coefficient is somewhat arbitrary and it depends on the value of D/T, that defines the “shock width""."," The value of the proportionality coefficient is somewhat arbitrary and it depends on the value of $\G/\Gu$ that defines the “shock width""."
+" We chose I'/T,,=0.9 obtaining n,orz’£z0.011? (the value of thelogarithmic factor is 5—10 for Γι=2— 10).", We chose $\G/\Gu=0.9$ obtaining $n_u \sigma_T z' \approx 0.01 \G_u^2$ (the value of thelogarithmic factor is $5-10$ for $\G_u =2-10$ ).
+" Choosing a different value of ['/T,,, to define the shock width, changes the coefficient, which in turn has a very weak effect on equation δ.."," Choosing a different value of $\G/\Gu$, to define the shock width, changes the coefficient, which in turn has a very weak effect on equation \ref{eq gmaxi}."
+" Now, since the length scale in the upstream frame, z, is shorter by a factor Γι than in the shock frame (i.e., 2’= Tz), and defining τι= n,oyz, we approximate the shock width"," Now, since the length scale in the upstream frame, z, is shorter by a factor $\G_u$ than in the shock frame (i.e., $z'=\Gu z$ ), and defining $\tau_u = n_u \sigma_T z$ , we approximate the shock width"
+core binary fraction from the primordial value.,core binary fraction from the primordial value.
+ Thus we do not agree with Ivanova et al. (, Thus we do not agree with Ivanova et al. (
+2005) that the binary fraction in the core will be depleted in time.,2005) that the binary fraction in the core will be depleted in time.
+ We also do not agree that models of elobular cluster evolution need necessarily include large populations of primordial binaries., We also do not agree that models of globular cluster evolution need necessarily include large populations of primordial binaries.
+ Our simulations have shown that the binary population in the core of a cluster is continually being replenished by stars from outside the core. many of which were previously in (he core.," Our simulations have shown that the binary population in the core of a cluster is continually being replenished by stars from outside the core, many of which were previously in the core."
+ This is a process we have termedconvection., This is a process we have termed.
+ We also find that the binary content of an evolved star cluster is dominated by exchange binaries provided that the stellar density is relatively high., We also find that the binary content of an evolved star cluster is dominated by exchange binaries provided that the stellar density is relatively high.
+ This is (rue of our moderate-size globular cluster models aud we expect it (o be true in more massive clusters., This is true of our moderate-size globular cluster models and we expect it to be true in more massive clusters.
+ We also show that increasing the primordial binary fraction does not necessarily lead (ο an increase in the final binary fraction in [acl 1 gives more scope for binary depletion., We also show that increasing the primordial binary fraction does not necessarily lead to an increase in the final binary fraction – in fact it gives more scope for binary depletion.
+ A kev ancl paracloxical result is that a final binary fraction Chat can be achieved by choosing a higher primordial binary fraction may also be replicated by choosing an initially lower binary fraction., A key and paradoxical result is that a final binary fraction that can be achieved by choosing a higher primordial binary fraction may also be replicated by choosing an initially lower binary fraction.
+ We find that the overall binary fraction of a cluster does not vary appreciably from the prinordial value as a cluster evolves., We find that the overall binary fraction of a cluster does not vary appreciably from the primordial value as a cluster evolves.
+ This is a result of binary destruction being balanced bv a greater rate of escape of single stars compared (o. binaries., This is a result of binary destruction being balanced by a greater rate of escape of single stars compared to binaries.
+ We also find that the prinordial binary frequency of a cluster is well preserved. outside of the cluster hall-mass radius., We also find that the primordial binary frequency of a cluster is well preserved outside of the cluster half-mass radius.
+ Therefore. observations of the current binary [fraction in these regions is a good indieator of the primordial binary fraction while determination of the core binary. lraction provides an upper limit.," Therefore, observations of the current binary fraction in these regions is a good indicator of the primordial binary fraction while determination of the core binary fraction provides an upper limit."
+ We acknowledge the generous support of the Cordelia Corporation and that of Edw Norton which has enabled AMNII to purchase GRAPE-G boards and supporting hardware., We acknowledge the generous support of the Cordelia Corporation and that of Edward Norton which has enabled AMNH to purchase GRAPE-6 boards and supporting hardware.
+ We (hank the anonymous referee for extremely helpful comments and especially for alerting us (o (he scaling considerations., We thank the anonymous referee for extremely helpful comments and especially for alerting us to the scaling considerations.
+"flux related to the SDSS automatic procedures, we visually inspected the 94 SDSS spectra above and found 11 objects in which the listed detection of [Ne V] emission is seriously affected by instrumental features and sky residuals.","flux related to the SDSS automatic procedures, we visually inspected the 94 SDSS spectra above and found 11 objects in which the listed detection of [Ne V] emission is seriously affected by instrumental features and sky residuals."
+ These objects were then removed from the sample., These objects were then removed from the sample.
+ We also noticed that the line parameters listed in the SDSS spectroscopic tables are not accurate for weak emission lines over a steep continuum., We also noticed that the line parameters listed in the SDSS spectroscopic tables are not accurate for weak emission lines over a steep continuum.
+" For instance, in about one quarter of SDSS QSOs, the [O ΠΙ3727 emission line (see the Discussion) appears as an absorption line due to an overestimate of the continuum."," For instance, in about one quarter of SDSS QSOs, the [O II]3727 emission line (see the Discussion) appears as an absorption line due to an overestimate of the continuum."
+" We then retrieved the SDSS spectra and performed a Gaussian fit to the emission lines we were interested in, deriving fluxes, luminosities, and rest frame equivalent widths."," We then retrieved the SDSS spectra and performed a Gaussian fit to the emission lines we were interested in, deriving fluxes, luminosities, and rest frame equivalent widths."
+ All the emission line parameters used in this work for SDSS objects are derived from our direct fits., All the emission line parameters used in this work for SDSS objects are derived from our direct fits.
+" Whenever we compared our line measurements with those performed by other Authors on the same SDSS spectra (e.g. for a subsample of type-2 QSOs in 2)), we found an excellent agreement."," Whenever we compared our line measurements with those performed by other Authors on the same SDSS spectra (e.g. for a subsample of type-2 QSOs in ), we found an excellent agreement."
+ The X/NeV vs Ny relation for the 83 SDSS QSOs selected above is shown in Fig. 2.., The X/NeV vs $N_H$ relation for the 83 SDSS QSOs selected above is shown in Fig. \ref{xnevsdss}.
+ Most objects are unobscured and populate the same region of local Seyfert 1 galaxies., Most objects are unobscured and populate the same region of local Seyfert 1 galaxies.
+" The median X/NeV ratio for SDSS QSOs is 370, almost identical to the value of 400 found for local Seyfert 1s."," The median X/NeV ratio for SDSS QSOs is 370, almost identical to the value of 400 found for local Seyfert 1s."
+" The few objects showing significant intrinsic X-ray absorption appear in the optical as intermediate type Seyferts (i.e. 1.8-1.9) rather than pure type-1 To check for possible outliers, e.g. blue SDSS QSOs with very low X/NeV ratios («15, as observed for local CT Seyferts), we also considered those objects in the Υ09 catalog showing significant [Ne V] emission and which were either"," The few objects showing significant intrinsic X-ray absorption appear in the optical as intermediate type Seyferts (i.e. 1.8-1.9) rather than pure type-1 To check for possible outliers, e.g. blue SDSS QSOs with very low X/NeV ratios $<15$, as observed for local CT Seyferts), we also considered those objects in the Y09 catalog showing significant [Ne V] emission and which were either"
+uodeliudepeudenut approach which does not assmue any articular shape for the profile.,model-independent approach which does not assume any particular shape for the profile.
+ While the uncertainties are larec. the flewre indicates a clear tendency of an Increasing 2 with radius.," While the uncertainties are large, the figure indicates a clear tendency of an increasing $\beta$ with radius."
+" The o, profile. as obtained from he Jeaus equation (eq. 2))."," The $\sigma_r$ profile, as obtained from the Jeans equation (eq. \ref{Jeans equation}) ),"
+ is shown in figure 7.., is shown in figure \ref{sigma_r profile}.
+ The analytical expression we used for ον eq. (7)).," The analytical expression we used for $\beta$, eq. \ref{beta analytic
+expression}) ),"
+ constrains to equal unity at large radii. corresponding o purely radial orbits.," constrains $\beta$ to equal unity at large radii, corresponding to purely radial orbits."
+ D99 analyzed N-body simulatious and derived values somewhat simaller than 1 at their Huitine radius. Grogy.," D99 analyzed N-body simulations and derived values somewhat smaller than 1 at their limiting radius, $6r_{200}$."
+ To allow for deviation frou unitv we also fit⋅ the expression⋅ (€.4|OoPEC. where a ds a free parameter which governs the asviuptotic vchavior of 9.," To allow for deviation from unity we also fit the expression $(C+a)\frac{(r/r_c)^2}{(r/r_c)^2+1}-C$, where $a$ is a free parameter which governs the asymptotic behavior of $\beta$."
+ The value of &=1 eave the best fit. also consistent with the value obtained by the independent mcthod at the largest radial point.," The value of $a=1$ gave the best fit, also consistent with the value obtained by the model-independent method at the largest radial point."
+ In l| we used the mass profile as derived from N-rav and lensing data together with the new data ou the projected velocity dispersion and the galaxy surface munber deusitv to derive the 3D profiles of the velocity anisotropy aud ealaxy nuuniber density., In \ref{The galaxy dynamical properties} we used the mass profile as derived from X-ray and lensing data together with the new data on the projected velocity dispersion and the galaxy surface number density to derive the 3D profiles of the velocity anisotropy and galaxy number density.
+ Alternatively. there are at least two different wavs to derive the total mass profile of the cluster directly from the data sets on ealaxy dynamics.," Alternatively, there are at least two different ways to derive the total mass profile of the cluster directly from the data sets on galaxy dynamics."
+ Perhaps the simplest approach is to use oulv the velocity caustics derived above. the auiplitude of which is related to the escape velocity. which is a tracer of the cluster mass profile.," Perhaps the simplest approach is to use only the velocity caustics derived above, the amplitude of which is related to the escape velocity, which is a tracer of the cluster mass profile."
+ The second approach is to use both data sets (of the projected velocity dispersion and ealaxy surface number deusitv) together with the Jeaus equation to fit an NEW inass profile., The second approach is to use both data sets (of the projected velocity dispersion and galaxy surface number density) together with the Jeans equation to fit an NFW mass profile.
+ Thus. we have two indepeucent methods for estimating the cluster lass profile.," Thus, we have two independent methods for estimating the cluster mass profile."
+ As meutioned already. D99 has shown that the 3D mass profile can be fairly well estimated directly frou," As mentioned already, D99 has shown that the 3D mass profile can be fairly well estimated directly from"
+"model with cosmological constant O4=0.7. Oy,=0.3. and a baryon density O1,=0.04. the Hubble constant is Z4,=100 hb km ! |. with h=0.7. and oy=0.8.","model with cosmological constant $\Omega_\Lambda= 0.7$ , $\Omega_\textrm{m} = 0.3$, and a baryon density $\Omega_\textrm{b} = 0.04$, the Hubble constant is $H_0=100$ h km $^{-1}$ $^{-1}$ , with $\textrm{h} = 0.7$, and $\sigma_8 = 0.8$."
+ The gravitational softening was set as ep;—7.5 ! kpe.," The Plummer-equivalent gravitational softening was set as $\epsilon_{Pl}=
+7.5$ $^{-1}$ kpc."
+ The code used to perform the simulation is the TREESPII code GADGET-2 Springeletal.(2001):(2005).," The code used to perform the simulation is the TREESPH code GADGET-2 \cite{Springel01, Springel05}."
+. The siniulation follows the evolution of 480* dark matter (DM) particles and as many barvonic eas particles [from redshift 2=49 (o -0., The simulation follows the evolution of $480^3$ dark matter (DM) particles and as many baryonic gas particles from redshift $z = 49$ to $z=0$.
+" The box for the simulation is a cube of side 192 ! Mpe. the DM and gas particles have initial masses ni;=4.62xLO’ IM. and ni,=6.03xLOS tM. respectively."," The box for the simulation is a cube of side 192 $^{-1}$ Mpc, the DM and gas particles have initial masses $\textrm{m}_{\textrm{\tiny DM}} = 4.62 
+\times 10^9$ $^{-1}$ $_\odot$ and $\textrm{m}_{\textrm{\tiny gas}} = 
+6.93 \times 10^8$ $^{-1}$ $_\odot$ respectively."
+ The physical processes involved in the simulation are eravitv. non-radiative hvdrodynamies. star formation. feedback from SNe with the effect of weak galactic outflows. radiative gas cooling and heating bv a uniform. ünme-dependent. »xhotoionizinge ultraviolet backeround.," The physical processes involved in the simulation are gravity, non-radiative hydrodynamics, star formation, feedback from SNe with the effect of weak galactic outflows, radiative gas cooling and heating by a uniform, time-dependent, photoionizing ultraviolet background."
+e The treatment of radiative cooling assumes an optically thin gas in CIE and uses only the primordial abunelances (hydrogen mass fraction N=0.76. helium Y=0.24).," The treatment of radiative cooling assumes an optically thin gas in CIE and uses only the primordial abundances (hydrogen mass fraction $=0.76$ , helium $=0.24$ )."
+ Metals generated bv the simulation itself are not considered for radiative cooling., Metals generated by the simulation itself are not considered for radiative cooling.
+5 A uniform. time-dependent UV background Laarcdt&Macau(1996) reionizes the Universe al z~6.," A uniform, time-dependent UV background \cite{Haardt96} reionizes the Universe at $z\sim6$."
+ Star formation is introduced following a hybrid multiphase model for the interstellar medium Springel&llernquist(2003)., Star formation is introduced following a hybrid multiphase model for the interstellar medium \cite{Springel03}.
+. The interstellar medium. where star formation Lakes place. is represented as cold clouds (cold gas) embedded in a hot gas.," The interstellar medium, where star formation takes place, is represented as cold clouds (cold gas) embedded in a hot gas."
+ Clouds are not accounted for individually in a given star-lormine particle. but rather μον are (treated. all together as a fraction of the total mass of the given star-forming particle.," Clouds are not accounted for individually in a given star-forming particle, but rather they are treated all together as a fraction of the total mass of the given star-forming particle."
+ Every gas particle is considered as composed of (wo parts. the hot gas. with its own mass and density. and the cold cloud fraction. temperature and density. determine (he relative abundances of the two components.," Every gas particle is considered as composed of two parts, the hot gas, with its own mass and density, and the cold cloud fraction, temperature and density determine the relative abundances of the two components."
+ Whenever star lormation takes place. a new star particle is spawnecl (with just a fraction of the starting gas particle). Uus increasing the number of star particles.," Whenever star formation takes place, a new star particle is spawned (with just a fraction of the starting gas particle), thus increasing the number of star particles."
+ Stars are created. following a Salpeter initial mass function Salpeter(1955).. and instantly produce metals and release energv as supernovae.," Stars are created, following a Salpeter initial mass function \cite{Salpeter55}, and instantly produce metals and release energy as supernovae."
+ Metals aid energy. are carried to the intracluster and intergalactic medium by galactic winds., Metals and energy are carried to the intracluster and intergalactic medium by galactic winds.
+ Winds are introduced in the simulation as mass outflows wilh rate equal to (vice the star formation rate and wilh a wind velocity of 360 kms !., Winds are introduced in the simulation as mass outflows with rate equal to twice the star formation rate and with a wind velocity of $360$ km $^{-1}$.
+ The output of the simulation consists of 102 boxes. equally spaced in the logarithm of ihe expansion [actor between z=9 and z=0.," The output of the simulation consists of 102 boxes, equally spaced in the logarithm of the expansion factor between $z=9$ and $z=0$."
+ As shown in previous work(e.g.. (2006)))," As shown in previous work(e.g., \cite{Ursino06}) )"
+ the X-ray emission above redshift 2 is negligible., the X-ray emission above redshift 2 is negligible.
+ For our work we therefore onlv used (he simulationup to redshilt 2., For our work we therefore only used the simulationup to redshift 2.
+ Desides (he improvement in spatial resolution compared to ourprevious work. (his simulation suites well our needs due to ils large scale that. allows us enough statistics for distant regions.," Besides the improvement in spatial resolution compared to ourprevious work, this simulation suites well our needs due to its large scale that allows us enough statistics for distant regions."
+More recent cosmological simulations have box sizes of al mostLOO | Mpe,More recent cosmological simulations have box sizes of at most100 $^{-1}$ Mpc
+The black-hole spin in AIB3 7 has been very accurately measured to he in climensiouless spin paraiuceter (Liuetal.2008).,The black-hole spin in M33 $-$ 7 has been very accurately measured to be in dimensionless spin parameter \citep{Liu08}.
+. The authors of this paper show that holes” (Shatecet aud MeCliutocketal. (2006)))., The authors of this paper show that \citet{Sha06} and \citet{McC06}) ).
+ However. etal(2003) noted that their spin derivation is model aud subject to possible svstematic errors.," However, \citet{Liu08} noted that their spin derivation is model-dependent and subject to possible systematic errors."
+ Leeetal.(2002) predicted the spin parameters of Nova Sco (X-ray Nova Scorpii 1991) aud Il Lupi (1U I7) to be ~OLS. with small effects after they were born in the explosion frou mass accretion: 1.6.. predicted them as natal.," \citet{Lee02} predicted the spin parameters of Nova Sco (X-ray Nova Scorpii 1994) and Il Lupi (4U $-$ 47) to be $\sim0.8$, with small effects after they were born in the explosion from mass accretion; i.e., predicted them as natal."
+ However Brownctal.(2007) showed that the rotational energv im such binaries scaled inversely with the donor mass at the time of conunonu-envelope evolution preceding the explosion in which the black hole was born., However \citet{BLMM07} showed that the rotational energy in such binaries scaled inversely with the donor mass at the time of common-envelope evolution preceding the explosion in which the black hole was born.
+ The donor masses of Nova Sco aud Ib Lupi are ~27M. whereasthat of M33 7 was ~SOM... so Brownctal.(2007). sugeested that the 3.15 day period of N33. 7 resulted. from a dark explosion: the high spin parameter would have resulted from mass accretion.," The donor masses of Nova Sco and Il Lupi are $\sim2\msun$ whereasthat of M33 $-$ 7 was $\sim80\msun$, so \citet{BLMM07} suggested that the $3.45$ day period of M33 $-$ 7 resulted from a dark explosion; the high spin parameter would have resulted from mass accretion."
+ This mass accretion would have had to take placeat wperciitical rate as we discuss in this letter., This mass accretion would have had to take placeat hypercritical rate as we discuss in this letter.
+ We breflv comunent on the lywpercritical accretion ov M/Mrg;g;~lO or groater (Brown&Weineart-rer1991)., We briefly comment on the hypercritical accretion for $\dot{M}/\dot{M}_{Edd}\sim10^3$ or greater \citep{Bro94}.
+ The scenario begius with Boudi accretion hrough the sonic poiut. which is often ereatlv larger han accretion at the Eddiugton lait.," The scenario begins with Bondi accretion through the sonic point, which is often greatly larger than accretion at the Eddington limit."
+ Because it had recone worked out for a value of 0.31«10!ALEay by Brown&Weineartuer(1991). aud because this value is in the widdle of those we shall use in stellar evolution. we use this value. although it could be much. ercater.," Because it had been worked out for a value of $0.31\times10^4\dot{M}_{Edd}$ by \citet{Bro94} and because this value is in the middle of those we shall use in stellar evolution, we use this value, although it could be much greater."
+ The Brown&Weinegartuer(1991) work had been carried out earlier in all detail bv IIouck&Chevalier(1991). and checked by Chevalier(1995)., The \citet{Bro94} work had been carried out earlier in all detail by \citet{Hou91} and checked by \citet{Che95}.
+.. We uote that there is still considerable controversy in the astroplivsical community about whether hypererifical accretion cau take place or iof., We note that there is still considerable controversy in the astrophysical community about whether hypercritical accretion can take place or not.
+" Tlowever. the general point we address is that if AL exceeds τμ, then some of the accretion energv uust be removed by meaus other than photons."," However, the general point we address is that if $\dot{M}$ exceeds $\dot{M}_{Edd}$, then some of the accretion energy must be removed by means other than photons."
+ In the case of lvpercritical accretion. this excess cucrey cau ve carried off bv neutrino pairs (Brown&Weingartucr199 D.," In the case of hypercritical accretion, this excess energy can be carried off by neutrino pairs \citep{Bro94}."
+. In the case of a neutron star. neutrino losses allow he matter flow to join smoothly onto the neutron star surface.," In the case of a neutron star, neutrino losses allow the matter flow to join smoothly onto the neutron star surface."
+ In the case of a black hole. the neutring losses et the matter flow smoothlv over the event horizon aud disappear iuto the black hole.," In the case of a black hole, the neutrino losses let the matter flow smoothly over the event horizon and disappear into the black hole."
+" Iu the work of Podsiadlowskietal. (2003).. we uote wo possible stages where hypercritieal (M/Mag2, 107)or supercritical acerction may take place."," In the work of \citet{Pod03}, , we note two possible stages where hypercritical $\dot M/\dot M_{Edd} \gsim 10^3$ )or supercritical accretion may take place."
+ Podsiadlowski evolve a binary with Ap;= 19211... Wsecondary=20M.. auc orbital period of 6.8 davs.," \citet{Pod03} evolve a binary with $M_{BH}=12\msun$ , $m_{secondary}=25\msun$ and orbital period of 6.8 days."
+photometry is accurate to 0.02 mags. and the velocity dispersions to about SU (see Bernardi et al.,"photometry is accurate to 0.02 mags, and the velocity dispersions to about $8$ (see Bernardi et al."
+ 2003a)., 2003a).
+ We work with de Vaucouleurs (1918) magnitudes and sizes. and SDSS colors throughout: these are colors measured within au aperture which scales with the de Vaucouleurs μαΠο radius iu the k-baud.," We work with de Vaucouleurs (1948) magnitudes and sizes, and SDSS colors throughout: these are colors measured within an aperture which scales with the de Vaucouleurs half-light radius in the $r$ -band."
+ The most important reason for analyzing a new sample (other than size) is that the old photometric reductions output bv the SDSS pipeline were incorrect (see documentation on DR2. the Second Data Release).," The most important reason for analyzing a new sample (other than size) is that the old photometric reductions output by the SDSS pipeline were incorrect (see documentation on DR2, the Second Data Release)."
+ A comparison of the properties of objects for which old aud new photometric reductions are available shows that the new corrected photometry has made most magnitudes slightly fainter (~0.13 mags). and most halt-light radii smaller (~ )).," A comparison of the properties of objects for which old and new photometric reductions are available shows that the new corrected photometry has made most magnitudes slightly fainter $\sim 0.13$ mags), and most half-light radii smaller $\sim$ )."
+ It is customary to report velocity dispersions at sole fraction (woe use 1/8) of the halflieht radius., It is customary to report velocity dispersions at some fraction (we use 1/8) of the half-light radius.
+ The half-lieht radii of the galaxies which cuter our sample are typically about 2 arcsecs. approximately iudepenudoeut of redshift. whereas the SDSS fiber used to measure the spectruni from which the velocity dispersion is estimated has a diameter of 3 arcsec.," The half-light radii of the galaxies which enter our sample are typically about 2 arcsecs, approximately independent of redshift, whereas the SDSS fiber used to measure the spectrum from which the velocity dispersion is estimated has a diameter of 3 arcsec."
+ Since the velocity dispersious of carly-type galaxies are kuown to merease towards the center (following Joregcusen et al., Since the velocity dispersions of early-type galaxies are known to increase towards the center (following rgensen et al.
+ 1995 we assume the sealing is X(rr)991 where ris the hal£ligbt radius). all ucasured velocity dispersions are aperture corrected’. aud it is these which are usually reported (the mean correction is )).," 1995 we assume the scaling is $\propto (r/r_e)^{0.04}$, where $r_e$ is the half-light radius), all measured velocity dispersions are `aperture corrected', and it is these which are usually reported (the mean correction is )."
+ Although the measured velocity dispersions have rot changed. the new photometry has chaneed the haltight radius.," Although the measured velocity dispersions have not changed, the new photometry has changed the half-light radius."
+ Hence. aperture corrected velocity dispersious differ from those associated with the old plotometiy (1.6.. hose reported iu Bernardi et al.," Hence, aperture corrected velocity dispersions differ from those associated with the old photometry (i.e., those reported in Bernardi et al."
+ 2003a) by less than one vercent (the new values ave laveer by a factor of 1.184)., 2003a) by less than one percent (the new values are larger by a factor of $1.1^{0.04}$ ).
+ Bernardi et al. (, Bernardi et al. (
+2003b) showed that the huninosities in their sample (~9000 ealaxies with :< 0.3) evolve: Aff.)=ALO)O52. where AL(2) denotes the mean absolute magnitude at redshift +.,"2003b) showed that the luminosities in their sample $\sim 9000$ galaxies with $z\le 0.3$ ) evolve: $M_*(z) = M_*(0) - 0.85z$, where $M_*(z)$ denotes the mean absolute magnitude at redshift $z$."
+" We fud similar evolution in the new sample: the main differcuce is that the new photometric reductions makeA, fainterby about 0.125 mags.", We find similar evolution in the new sample: the main difference is that the new photometric reductions make $M_*$ fainter by about 0.125 mags.
+ The evolution is consistent with that of a passively aging »pulatiou., The evolution is consistent with that of a passively aging population.
+ Thus. in the analysis which follows. we will © careful to separate trends with redshift that are due o the magnitude lait of the sample. from trends that απο due to evolution.," Thus, in the analysis which follows, we will be careful to separate trends with redshift that are due to the magnitude limit of the sample, from trends that are due to evolution."
+ Maux of the following figures preseut ueasureineuts made dm a series of narrow redshift bins: L02xi«cQT. OOFxoiozQ(LQ09. 0009.—loxO2. A12x:<O15. and 0.15xs<0.2.," Many of the following figures present measurements made in a series of narrow redshift bins: $0.02\le z<0.07$, $0.07\le z<0.09$, $0.09\le z < 0.12$, $0.12\le z< 0.15$, and $0.15\le z<0.2$."
+" In all of these Ποπος, sviubols show the median value of the variable ou hne v-axis as a function of the observable ou the x-axis. error bars show the rius error of the median. and hashed regious indicate where the central of the sample in cach redshift bin lies."," In all of these figures, symbols show the median value of the variable on the y-axis as a function of the observable on the x-axis, error bars show the rms error of the median, and hashed regions indicate where the central of the sample in each redshift bin lies."
+" Figure ld displavs correlations between Πορτν, size [ει aud velocity dispersion σ in this sample. ("," Figure \ref{lvrm} displays correlations between luminosity, size $R_r$, and velocity dispersion $\sigma$ in this sample. ("
+"The light radius in kpc. &,. is defined similarly to the quantity called Π by Bernardi et al.","The half-light radius in kpc, $R_r$, is defined similarly to the quantity called $R_o$ by Bernardi et al."
+ 2003a. for the r-baud.)," 2003a, for the $r$ -band.)"
+ The dashed lines show fits to these correlations. stmunarizec in Table 1..," The dashed lines show fits to these correlations, summarized in Table \ref{MLcov}."
+ Tn all cases. luminosities have been corrected or evolution by 0.852.," In all cases, luminosities have been corrected for evolution by $z$."
+ In the Table. AZ. aud σα are the uean and rms values of a Gaussian fit to the evolution-corrected ας huuinosity function. Q quantifies the rate of Iuniuositv evolution. aud R. aud V are the values of ogyy(size/h type) aud logyy(velocity dispersion/liu 1) or an L. galaxy.," In the Table, $M_*$ and $\sigma_M$ are the mean and rms values of a Gaussian fit to the evolution-corrected $r$ -band luminosity function, $Q$ quantifies the rate of luminosity evolution, and $R_*$ and $V_*$ are the values of $\log_{10}$ $h^{-1}$ kpc) and $\log_{10}$ (velocity dispersion/km $^{-1}$ ) for an $L_*$ galaxy."
+ Iu general. these values are all qualitatively simular to those reported in Bernardi et al. (," In general, these values are all qualitatively similar to those reported in Bernardi et al. ("
+20035). but here are quantitative differences. most noticably iu the uean size Πο. which is simaller by about ten percent.,"2003b), but there are quantitative differences, most noticably in the mean size $R_*$, which is smaller by about ten percent."
+ Table 1 also reports values of a ΠΙΟ of pairwise correlations in the following format., Table \ref{MLcov} also reports values of a number of pairwise correlations in the following format.
+ Caven a pair of observables X aud Y. (YVAINfoyy=&£xyv(X.Noy.," Given a pair of observables $X$ and $Y$, $\langle Y-Y_*|X\rangle/\sigma_{YY} = \xi_{XY}(X-X_*)/\sigma_{XX}$."
+ The results of a similar analysis of the correlations between color. magnitude aud velocity dispersion is given in Table 2..," The results of a similar analysis of the correlations between color, magnitude and velocity dispersion is given in Table \ref{MLcmag}."
+ Figure 2 shows the coloranagnitude relation im our sample., Figure \ref{cmag} shows the color-magnitude relation in our sample.
+ The top of seriespancls shows the raw mcasureient. and bottom panels show the result of accounting for evolution by adding 0.852 mags to the maguitudes. aud adding 0.3: mages to the colors (see Table 2)).," The top series of panels shows the raw measurement, and bottom panels show the result of accounting for evolution by adding $0.85z$ mags to the magnitudes, and adding $0.3z$ mags to the colors (see Table \ref{MLcmag}) )."
+ Ou both tiers. the left-ost paucl shows the color-magnitude relation in the full sample. aud the other two panels show the relation in siunall bius in velocity dispersion.," On both tiers, the left-most panel shows the color-magnitude relation in the full sample, and the other two panels show the relation in small bins in velocity dispersion."
+ The füsgure makes two miportaut poiuts: (1) at fixed velocity dispersion. there is no relation between color aud magnitude. aud (i) the slope of the relation in the loftinost panel is approximately the same im all redshift bius.," The figure makes two important points: (i) at fixed velocity dispersion, there is no relation between color and magnitude, and (ii) the slope of the relation in the leftmost panel is approximately the same in all redshift bins."
+ We also find no evidence that the scatter around the mean relations changes with redshift. but placing a precise limit of the evolution is difficult. because the expected sienal is small so the auswer is scusitive to uncertainties in the SDSS photometry (0.02 1iags) aud the &-correction.," We also find no evidence that the scatter around the mean relations changes with redshift, but placing a precise limit of the evolution is difficult, because the expected signal is small, so the answer is sensitive to uncertainties in the SDSS photometry (0.02 mags) and the $k$ -correction."
+Compton scatterime (IC) with thermal photons im the Wlein-Nishina regiue. with a cooling time-scale τι~3.108(eui/80TCV)l(es/2keV) d. which is shorter than the piondecay time Tey22.6«105s5.=5.5«10s(eti/30 TeV).,"Compton scattering (IC) with thermal photons in the Klein-Nishina regime, with a cooling time-scale $\tau_{\rm IC} \sim 3 \times 10^{-4}~{\rm s}~
+(\epsilon_{\pi^{+}}/30{\rm TeV})^{-1} (\epsilon_\gamma/2{\rm keV})^{-4}$ , which is shorter than the piondecay time $\tau_{decay} \simeq
+2.6 \times 10^{-8}~{\rm s}~\gamma_{\pi^{+}}=5.5\times 10^{-3}~{\rm s}~
+(\epsilon_{\pi^{+}}/30~{\rm TeV})$ ."
+ Thus the pious undereo some IC cooling before they decay. aud the typical neutrino enerey is down by a factor f.=(0.3/5.5)?~0.23 relative to the ἐν," Thus the pions undergo some IC cooling before they decay, and the typical neutrino energy is down by a factor $f_{c}=(0.3/5.5)^{1/2}\sim 0.23$ relative to the $\epsilon_{\nu,th}$."
+ Therefore for discrete sources. the typical neutrine euergv is The final πιο huninositv is = Logf. y? plu )..," Therefore for discrete sources, the typical neutrino energy is The final neutrino luminosity is = f_c )^2 )^4."
+"3) Since the spiudown buninositv is Le,=LA«LO?tore:SDRPSs) νο the neutro eiission efficiency is then 0.01 RAED) Lb We"," Since the spindown luminosity is $L_{sd}=1.5\times 10^{34}
+{\rm erg/s}B_{p,15}^2 R_6^6 (P/5{\rm s})^{-4}$ , the neutrino emission efficiency is then = 0.04 )^2 ) R_6^4 )^4."
+" (Gs)asstune that this Iuninuositv is beamed iuto a sweep-averaged solid angle AQ,~0.1. which is typical for a polar cap angle 0.01 and a moderate inclination angele of the rotator."," We assume that this luminosity is beamed into a sweep-averaged solid angle $\Delta \Omega_\nu \sim 0.1$, which is typical for a polar cap angle $\sim 0.01$ and a moderate inclination angle of the rotator."
+" A smaller/lareer AO, increasesdecreases the on-heam neutrino flux. but decreases/iucreascs the probability of on-heam detection."," A smaller/larger $\Delta \Omega_{\nu}$ increases/decreases the on-beam neutrino flux, but decreases/increases the probability of on-beam detection."
+ For au ou-beam observer. the neutrino nuniber fix at earth is MGT)? Bas «(Enteste temennam where D is the distance to the (psource.," For an on-beam observer, the neutrino number flux at earth is = )^2 ) ^3 )^4, where $D$ is the distance to the source."
+" The probability of detecting a neutrinoduduced upward muon with planned netitrino telescopes is Z2,wy&1.3«10(e,'TeV) (Halzeu Ilooper 2002). eiviug au on-beam upward muon event rate dAdtidiserotey~L.7(i028)144) REE)"," The probability of detecting a neutrino-induced upward muon with planned neutrino telescopes is $P_{\nu\to\mu}\simeq 1.3\times 10^{-6} (\epsilon_{\nu}/{\rm TeV})$ (Halzen Hooper 2002), giving an on-beam upward muon event rate ) 1.7 ) ^3 )^4."
+ The chances for the observer to be in the neutrino beau are not large., The chances for the observer to be in the neutrino beam are not large.
+ Nonetheless. there is a small but finite probability for directly detecting some neutrinos frou these objects.," Nonetheless, there is a small but finite probability for directly detecting some neutrinos from these objects."
+ In Table 1. we give the predicted muon event rates for the four magnuetar candidates which may be neutrino loud under favorable conditions. assumnius an ou-bea observation.," In Table 1, we give the predicted muon event rates for the four magnetar candidates which may be neutrino loud under favorable conditions, assuming an on-beam observation."
+" Since other mnaguetars all lic below the most favorable f,=1/2 death valley. they are definetely below the photomesou threshold. aud we do uot consider them as neutrino cussion candidates."," Since other magnetars all lie below the most favorable $f_g=1/2$ death valley, they are definetely below the photomeson threshold, and we do not consider them as neutrino emission candidates."
+ Frou Table 1. we see that SCR 1900111 and TE 10185-5937 may be detected by: kan? telescopes with several years of operation. if they are above the photomesonu threshold aud if their neutrino beams sweep the Earth.," From Table 1, we see that SGR 1900+14 and 1E 1048-5937 may be detected by $^3$ telescopes with several years of operation, if they are above the photomeson threshold and if their neutrino beams sweep the Earth."
+" Accoupauving the neutrinos there should also be clectromaguctic signals frou z""-decay aud z! cooling aud cascading.", Accompanying the neutrinos there should also be electromagnetic signals from $\pi^{0}$ -decay and $\pi^{+}$ cooling and cascading.
+ Thehieh 5D aud 55 pair-formation aud photon splitting opacity in the strong maeguetic fields may deerade the typical photon energy to < 140 MeV. below the EGRET baud (c.c. Ibuxiug. Baring Couthier 1997: Baring Tarding 2001). but it mav fall into the INTEGRAL baud.," Thehigh $\gamma B$ and $\gamma\gamma$ pair-formation and photon splitting opacity in the strong magnetic fields may degrade the typical photon energy to $\lesssim$ 40 MeV, below the EGRET band (e.g. Harding, Baring Gonthier 1997; Baring Harding 2001), but it may fall into the INTEGRAL band."
+ A direct inference from the above proposal is that the eutire population of vouug magnetars in the universe will contribute to a diffuse neutrino backeround. before crossing the neutrino death valley.," A direct inference from the above proposal is that the entire population of young magnetars in the universe will contribute to a diffuse neutrino background, before crossing the neutrino death valley."
+ The number flux of this backeround cau be generally estimated as where Dy~1075 cin is the Hubble distance. aud εν is the typical energy of the ueutrino background.," The number flux of this background can be generally estimated as where $D_H \sim 10^{28}$ cm is the Hubble distance, and $\bar
+\epsilon_\nu$ is the typical energy of the neutrino background."
+ The inner mteeral is the average total neutriuo energv fiueuce oer luaeuctar enütted towards earth duriug its neutrino-oud Tite fine Taga75«105 sy. which is based on he known magnetars being marginal neutrino enütters.," The inner integral is the average total neutrino energy fluence per magnetar emitted towards earth during its neutrino-loud life time $\tau_{mag,\nu} \sim 5\times 10^3$ yr, which is based on the known magnetars being marginal neutrino emitters."
+ Since 9imaenetars have heen discovered in the Galaxy with ypical ages of 101 vr. the local (redslift 2=0) maguctar arth rate cau be conservatively estimated Ίο).~10?xvrtealaxyTR4792«107?vr!MpeOR s fora nunber density of ealaxies n4=0.02Mpc* (Allen 1973).," Since 9 magnetars have been discovered in the Galaxy with typical ages of $10^4$ yr, the local (redshift $z=0$ ) magnetar birth rate can be conservatively estimated ${\cal R}(0) \simeq 10^{-3}~{\rm yr^{-1} galaxy^{-1}} {\cal R}_{-3} 
+\simeq 2\times 10^{-5}~{\rm yr^{-1} Mpc^{-3}} {\cal R}_{-3}$ , for a number density of galaxies $n_g=0.02~{\rm Mpc}^{-3}$ (Allen 1973)."
+ Asuniug that the magnetar birth rate follows the star ornüug rate. R(:)=RIOD foro: «2 (Lilly et al.," Assuming that the magnetar birth rate follows the star forming rate, ${\cal R}(z) \simeq {\cal R}(0) (1+z)^3$ for $z<2$ (Lilly et al."
+ 1996)., 1996).
+ The time-dependent beaming paraicter faf) Qvhich is the fraction of maguctars whose neutrino beams are directed towards us. so the sweep-averaged solid angle of the neutrino beam is AO(f)= πμ) cancels out.," The time-dependent beaming parameter $f_b(t)$ (which is the fraction of magnetars whose neutrino beams are directed towards us, so the sweep-averaged solid angle of the neutrino beam is $\Delta
+\Omega(t)=4\pi f_b(t)$ ) cancels out."
+ The outer integral is over the Wubble volume., The outer integral is over the Hubble volume.
+ For remote maegnetars. the neutrino flux of au individual source drops as D7 while the total munber of maguctars increases as D? foy ;<1.," For remote magnetars, the neutrino flux of an individual source drops as $D^{-2}$ while the total number of magnetars increases as $D^3$ for $z\ll1$."
+ Therefore most of the diffuse neutiiuo Cluission cones from the farthest magnuetars whose birth rate is the highest., Therefore most of the diffuse neutrino emission comes from the farthest magnetars whose birth rate is the highest.
+" For wouug iaeuetars. the time-dependent neutrine ""numositv may be estimated as in 833."," For young magnetars, the time-dependent neutrino luminosity may be estimated as in 3."
+ There are some joficeable differences. however.," There are some noticeable differences, however."
+" For cxaimple. due to radiation reaction aud possible pair screcning effect. 5, 1."," For example, due to radiation reaction and possible pair screening effect, $\eta_p \ll 1$ ."
+ On the other haud. the pair screening altitude could be uuch higher than the altitude where pious are generated. so that pions could undergo substautial reacceleration fore decaviug.," On the other hand, the pair screening altitude could be much higher than the altitude where pions are generated, so that pions could undergo substantial reacceleration before decaying."
+ As a result f. could be c1., As a result $f_c$ could be $\gg 1$.
+ Notice hat these uncertainties oulv influence the typical energy of the neutrino background. ἐν. but do not iufluence he number counts of the neutro backeround (17)). which can be estimated as follows.," Notice that these uncertainties only influence the typical energy of the neutrino background, $\bar \epsilon_\nu$ , but do not influence the number counts of the neutrino background \ref{barphinu}) ), which can be estimated as follows."
+" The time-depeudoeut )cutrino luminosity is L,(f)=aA, (f)enzi(f)ez. where Aydt)1=ποΠο is the time-dependent polar. cap"," The time-dependent neutrino luminosity is $L_\nu (t) = A_{pc}(t) c n_{\pi^{+}}(t)\bar \epsilon_\nu$ , where $A_{pc}(t)=\pi \Omega(t)R^3/c$ is the time-dependent polar cap"
+Data reduction was performed with the MIRIAD package (Sault. Teuben Wright 1995).,"Data reduction was performed with the MIRIAD package (Sault, Teuben Wright 1995)."
+ In most cases naturally weighted. cleaned maps. Were used for image analysis.," In most cases naturally weighted, cleaned maps, were used for image analysis."
+ We were able to produce a combined map from all the data using visibilities calibrated individually at each epoch., We were able to produce a combined map from all the data using visibilities calibrated individually at each epoch.
+ Fig 2 presents a 4.8 GHz map of the field of XTE J1701-462 in which all sixteen datasets were combined. resulting in slightly more than 24 hr on-source.," Fig 2 presents a 4.8 GHz map of the field of XTE J1701-462 in which all sixteen datasets were combined, resulting in slightly more than 24 hr on-source."
+ The radio counterpart of XTE 11701-462 is clearly detected at RA 17:00:58.43 Dee -46:11:08.44. with an uncertainty of about 0.3 aresec in each coordinate.," The radio counterpart of XTE J1701-462 is clearly detected at RA 17:00:58.43 Dec -46:11:08.44, with an uncertainty of about 0.3 arcsec in each coordinate."
+ In order to compare the radio and X-ray properties of XTE we have made an estimate of which part of the Z the source was on at or close to the time of each radio observation., In order to compare the radio and X-ray properties of XTE we have made an estimate of which part of the 'Z' the source was on at or close to the time of each radio observation.
+ XTE 1701—162. was observed 97 times with RXTE between 2006 Jan 22 - March 20. and once on 2006 April 23. providing simultaneous X-ray coverage for three of our radio observations and X-ray covergage within a few hours for the others.," XTE J1701–462 was observed 97 times with RXTE between 2006 Jan 22 - March 20, and once on 2006 April 23, providing simultaneous X-ray coverage for three of our radio observations and X-ray covergage within a few hours for the others."
+ The RXTE observations before March 20 were all classified according to their iming properties and position in X-ray colour-colour diagrams by Homan et (007) and a separate analysis of the 2006 April 23 RXTE observation was done for this paper., The RXTE observations before March 20 were all classified according to their timing properties and position in X-ray colour-colour diagrams by Homan et (2007) and a separate analysis of the 2006 April 23 RXTE observation was done for this paper.
+ The location along he Z tracks was estimated for the RXTE observations that were closest in time to our radio observations. by determining their xosition (upper/middle/lower) along the branches of full Z tracks hat were traced out in various time intervals (see Fig.," The location along the Z tracks was estimated for the RXTE observations that were closest in time to our radio observations, by determining their position (upper/middle/lower) along the branches of full Z tracks that were traced out in various time intervals (see Fig."
+ 3 in Homan et 22007)., 3 in Homan et 2007).
+ These estimates are indicated in Table |., These estimates are indicated in Table 1.
+ The radio detections and upper limits are also indicated as a function of branch of the Z in Fig 2., The radio detections and upper limits are also indicated as a function of branch of the Z in Fig 2.
+ Our observations appear to indicate that XTE J1701-462 is most likely to be detected as a radio source when on the horizonta or normal branches (HB. NB) of the Z. There are no detections on the flaring branch. including the most stringent upper limi (observation P3.," Our observations appear to indicate that XTE J1701-462 is most likely to be detected as a radio source when on the horizontal or normal branches (HB, NB) of the Z. There are no detections on the flaring branch, including the most stringent upper limit (observation P)."
+ This is consistent with the behaviour noted firs noted for GX 17-2 by Penninx et al. (, This is consistent with the behaviour noted first noted for GX 17+2 by Penninx et al. (
+1988) and possibly universa or all Z sources (Penninx 1989).,1988) and possibly universal for all Z sources (Penninx 1989).
+ These periods of relatively strong radio emission almost certainly correspond to the formation ofjet- ike outflows and/or their interaction with the surrounding medium (as directly observed in the Z/Z-like sources Sco X-I1 and Cir Fomalont et al., These periods of relatively strong radio emission almost certainly correspond to the formation of jet-like outflows and/or their interaction with the surrounding medium (as directly observed in the Z/Z-like sources Sco X-1 and Cir X-1; Fomalont et al.
+ 2001: Fender et al., 2001; Fender et al.
+ 2004)., 2004).
+ In particular. based orimarily on results for Sco X-1. Migliari Fender (2006) suggest hat the HB / NB vertex may correspond to the point at which the most powerful eject events occur. and that the jet is suppressed on he FB.," In particular, based primarily on results for Sco X-1, Migliari Fender (2006) suggest that the HB / NB vertex may correspond to the point at which the most powerful eject events occur, and that the jet is suppressed on the FB."
+ The Z sources are likely to be accreting persistently (in most cases) at close to the Eddington limit., The Z sources are likely to be accreting persistently (in most cases) at close to the Eddington limit.
+ It may be useful to compare hem to black hole X-ray binaries accreting at comparably high Eddington ratios. such as GRS 19154105.," It may be useful to compare them to black hole X-ray binaries accreting at comparably high Eddington ratios, such as GRS 1915+105."
+ Both classes of object make dramatic and rapid state transition and are associated with episodic production of powerful relativistic jets., Both classes of object make dramatic and rapid state transition and are associated with episodic production of powerful relativistic jets.
+ Such objects may. in turn. be our best ‘local’ equivalents of quasars accreting at very high rates at redshifts of ς=1.," Such objects may, in turn, be our best `local' equivalents of quasars accreting at very high rates at redshifts of $z \geq 1$."
+ Therefore careful comparison of objects such as XTE J1701-462 with black holes acereting at high rates may provide our best test of the effects of e.g. event horizons. static limits on accretion and jet formation.," Therefore careful comparison of objects such as XTE J1701-462 with black holes accreting at high rates may provide our best test of the effects of e.g. event horizons, static limits on accretion and jet formation."
+ The two-point radio spectra of these detections are plotted in Fig 3., The two-point radio spectra of these detections are plotted in Fig 3.
+" It is interesting to note that for the majority of the radio detections of XTE H701-462 the source radio spectrum is flat / inverted (spectral index à.=OQ. where S,x ov)."," It is interesting to note that for the majority of the radio detections of XTE J1701-462 the source radio spectrum is flat / inverted (spectral index $\alpha \geq 0$, where $_{\nu} \propto
+\nu^{\alpha}$ )."
+ This seems to be in contrast to most emission associated. with transient outbursts from X-ray binaries which generally has an optically thin spectrum (aX 1.6).," This seems to be in contrast to most emission associated with transient outbursts from X-ray binaries which generally has an optically thin spectrum $\alpha \leq
+-0.6$ )."
+ Optically thin emission is also observed in Cir X-| (Fender et al., Optically thin emission is also observed in Cir X-1 (Fender et al.
+ 1998: 2004) and Sco X-I (Fomalont et. al., 1998; 2004) and Sco X-1 (Fomalont et al.
+ 20013. although in both sources episodes of flat-spectrum core emission has been seen.," 2001), although in both sources episodes of flat-spectrum core emission has been seen."
+ The flat/inverted radio spectrum is in fact more reminiscent of the steady. flat-spectrum radio emission observed from black holes in hard X-ray states (Fender 2001). although Migliari Fender (2006) do suggest that it should also be observed in hard state neutron stars.," The flat/inverted radio spectrum is in fact more reminiscent of the steady, flat-spectrum radio emission observed from black holes in hard X-ray states (Fender 2001), although Migliari Fender (2006) do suggest that it should also be observed in hard state neutron stars."
+ Such flat-spectrum emission is believed to arise in a partially self-absorbed jet (Blandford Kónnigl 1979: Kaiser 2006)., Such flat-spectrum emission is believed to arise in a partially self-absorbed jet (Blandford Könnigl 1979; Kaiser 2006).
+ Given the rather poor coverage and, Given the rather poor coverage and
+do not proceed rapidly.,do not proceed rapidly.
+ Doiual anion destruction mechanisms are by reaction with H and and by photodetachiment., Dominant anion destruction mechanisms are by reaction with H and $^+$ and by photodetachment.
+ Photodetachinen rates were calculated according to Equation (2) of Milareal.(2007)., Photodetachment rates were calculated according to Equation (2) of \citet{mil07}.
+. As a result. of its large electron detachinent energv. t1ie photodetachinent rate thus calculated is ~100 times less than the value in the RATEOG database.," As a result of its large electron detachment energy, the $^-$ photodetachment rate thus calculated is $\sim100$ times less than the value in the RATE06 database."
+" Carbon chain anions ul η=2— 7). aud C,H (η=2. 1.6) have yee. shown to react witi atomic nitrogen and result in t formation of products that include the nitrile anions C,,N 0 (i0=1.3.5) (Eichelbergerοἱal. 2007)."," Carbon chain anions $_n^-$ $n=2-7$ ), and $_n$ $^-$ $n=2, 4, 6$ ) have been shown to react with atomic nitrogen and result in the formation of products that include the nitrile anions $_n$ $^-$ $n=1, 3, 5$ ) \citep{eic07}."
+. We have included these reactions in he model as part of our investigation iuto the yossible mechanistus for the formaion of auc the receuly ¢iscovered CaN (ThadcleIsetal. 2008)., We have included these reactions in the model as part of our investigation into the possible mechanisms for the formation of $^-$ and the recently discovered $_3$ $^-$ \citep{tha08}.
+. Branching ratios were calculated TOL EichelP)ergeretal.(2007) 5 ο‘ginal experimeuta data by Veronica Biebaum (private comumunication)., Branching ratios were calculated from \citet{eic07}' 's original experimental data by Veronica Bierbaum (private communication).
+" The braiching ratio or the associative electron detachrent (AED) »roduct. channel could not. be derive [rom he experimental data. so we have assuiued. arbitrarily, a ratio of 0.5."," The branching ratio for the associative electron detachment (AED) product channel could not be derived from the experimental data, so we have assumed, arbitrarily, a ratio of 0.5."
+ This assumption coustitutes potentially the most signilicaut source of error in these reaction rates., This assumption constitutes potentially the most significant source of error in these reaction rates.
+ Mutual neutralisation reactions between anions and catious were shown by Lepp&Dalearuo(1988) to have important ellects ou interstellar chemistry., Mutual neutralisation reactions between anions and cations were shown by \citet{lep88} to have important effects on interstellar chemistry.
+" Thus. reactions of the kinel ~3N4+Y(1)X* have been included for the twenty most abundant catious ). and for all auious ).in the model. with a rate coefficient of 75x10CD/300)"" cm? ! (seeHarada&Herbst2008.(or exzinple).."," Thus, reactions of the kind ^+ + }^- + have been included for the twenty most abundant cations $^+$ ), and for all anions $^-$ ), in the model, with a rate coefficient of $\times10^{-8}(T/300)^{-0.5}$ $^3$ $^{-1}$ \citep[see][for example]{har08}. ."
+ The final chemical network contains 126 e@as-plase species coupled by 5539 reactious., The final chemical network contains 426 gas-phase species coupled by 5539 reactions.
+ Initially the rate equations are solved as a funetion of radius (inthesamewayasMil-laretal.2¢)00).. starting [firon Lol? Cn and 1roving out (o ry=—20x1013 cin where photodissociation destroys all the molecules (apar [rom H»).," Initially the rate equations are solved as a function of radius \citep[in the same way as][]{mil00}, starting from $r_i=10^{15}$ cm and moving out to $r_f=3\times10^{18}$ cm where photodissociation destroys all the molecules (apart from $_2$ )."
+ Then in a separate routine the ¢relmical abuudauces iusile the cdesilv-enhatced shells are caleulated as a fuicllou of racius., Then in a separate routine the chemical abundances inside the density-enhanced shells are calculated as a function of radius.
+ The chemical rate equatio15 of a cleisltπι packet lying halfway betweeLa shells iner and outer radius fal a radius ry with «ensily (9+1)n(ry) are solved. starting [rou r; aud moving out to 2xLO” cnlw lere the density is sullicieutly low that 4he she uo longer makes aly sigIcanl c‘Outrilion to the total amount of matter il1 le 11odel.," The chemical rate equations of a density-enhanced packet lying halfway between a shell's inner and outer radius (at a radius $r_p$ with density $(\beta+1) n(r_p)$ ), are solved, starting from $r_i$ and moving out to $2\times10^{17}$ cm where the density is sufficiently low that the shell no longer makes any significant contribution to the total amount of matter in the model."
+ The motion of the ceusity-enhanced stells iu the outflow are Followed so lia they move oward over tiiJe. SVLironised with je outward radial motion ve deuse packet.," The motion of the density-enhanced shells in the outflow are followed so that they move outward over time, synchronised with the outward radial motion of the dense packet."
+ The adiation [field is culatec for the deuse pac‘ket at each tii1e- taking iuO account tle new positious | the stells.," The radiation field is recalculated for the dense packet at each time-step, taking into account the new positions of all the shells."
+ The clensity factor (9). he inter-sliel W.paciug and thes rell thickness are iclenical for every shell so that the chemical yttclances Caculated for the cleuse pacsel al a give1 radius represent the abunclauces in a densitv-euliauced. shell centered at that raclius.," The density factor $\beta$ ), the inter-shell spacing and the shell thickness are identical for every shell so that the chemical abundances calculated for the dense packet at a given radius represent the abundances in a density-enhanced shell centered at that radius."
+ The effects on the model results of variations in the mass-loss rate. radiation field strength. eas-lo-dust ratio. shell thickuess. inter-s1ell spaciiο ΠΠ factor aud stellar distauce have been analysed.," The effects on the model results of variations in the mass-loss rate, radiation field strength, gas-to-dust ratio, shell thickness, inter-shell spacing, density-enhancement factor and stellar distance have been analysed."
+ The resits presetted in this study. are forthe inodel tiat Ies the parameters that we believe best natch the observational COnlslrallls., The results presented in this study are forthe model that uses the parameters that we believe best match the observational constraints.
+ Moclification of these parameters. particularly thliose hat affect the radiation field strength.," Modification of these parameters, particularly those that affect the radiation field strength,"
+ , 
+surveys obtained through STScEs DSS Web interface.,surveys obtained through STScI's DSS Web interface.
+ All of the data analysis presented in this paper was carried out using the Interactive Data Language (IDL)., All of the data analysis presented in this paper was carried out using the Interactive Data Language (IDL).
+ We are greatly indebted to all who contributed to the various IDL User's Libraries: the routines of the IDL Astronomy Users Library (maintained by Wayne Landsman) have been used particularly extensively., We are greatly indebted to all who contributed to the various IDL User's Libraries; the routines of the IDL Astronomy User's Library (maintained by Wayne Landsman) have been used particularly extensively.
+ HE thanks the Time Allocation Committee of the Institute for Astronomy for their support of the spectroscopic follow-up observations of eBCS clusters., HE thanks the Time Allocation Committee of the Institute for Astronomy for their support of the spectroscopic follow-up observations of eBCS clusters.
+ HE also gratefully acknowledges partial financial support from a European Union EARA Fellowship. SAO contract SV4-64008. and NASA LTSA grant NAG 45-8253.," HE also gratefully acknowledges partial financial support from a European Union EARA Fellowship, SAO contract SV4-64008, and NASA LTSA grant NAG 5-8253."
+ ACE. ACF and CSC thank the Royal Society for support.," ACE, ACF and CSC thank the Royal Society for support."
+ This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service. provided by the NASA-Goddard Space Flight Center. and the NASA/IPAC Extragalactic Database (NED).," This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA-Goddard Space Flight Center, and the NASA/IPAC Extragalactic Database (NED)."
+medium-resolution optical wavelength spectra. spanning a wavelength range from 3474 to 7418A.,"medium-resolution optical wavelength spectra, spanning a wavelength range from 3474 to 7418."
+". All data were taken using the 3"" slit and a mmm! grating. producing a resolution of À/AA~ 1.000 at6.000A.."," All data were taken using the 3"" slit and a $^{-1}$ grating, producing a resolution of $\lambda / \Delta\lambda\sim$ 1,000 at."
+ Exposure times varied widely based on the magnitude of the star. but ranged from ~3 to 300 s. The two-dimesional spectra were bias-subtracted and flat-fielded by the observatory standard data reduction procedure.," Exposure times varied widely based on the magnitude of the star, but ranged from $\sim$ 3 to 300 s. The two-dimesional spectra were bias-subtracted and flat-fielded by the observatory standard data reduction procedure."
+ After extraction. and the one-dimensional spectra were wavelength calibrated.," After extraction, and the one-dimensional spectra were wavelength calibrated."
+ We used standard star spectra taken each night of the observations to produce sensitivity curves and flux calibrated spectra., We used standard star spectra taken each night of the observations to produce sensitivity curves and flux calibrated spectra.
+ An extinetion correction file from Kitt Peak National Observatory was also applied to correct for major atmospheric effects., An extinction correction file from Kitt Peak National Observatory was also applied to correct for major atmospheric effects.
+ The standard calibration. stars used to create the sensitivity curve are BD+26 2606. BD+33 2642. BD+17 4708. Feige 34. and HD 19445.," The standard calibration stars used to create the sensitivity curve are BD+26 2606, BD+33 2642, BD+17 4708, Feige 34, and HD 19445."
+ Flux calibration and extinetion corrections were performed using the package and the task in IRAF., Flux calibration and extinction corrections were performed using the package and the task in IRAF.
+ We used a combination of observatory logs and night-to-night comparison of the sensitivity function to determine whether the calibrated target spectra were reliable enough to use in color analysis., We used a combination of observatory logs and night-to-night comparison of the sensitivity function to determine whether the calibrated target spectra were reliable enough to use in color analysis.
+ We rejected some of the initial data sample based on concerns about non-photometric conditions and atmospheric distortion., We rejected some of the initial data sample based on concerns about non-photometric conditions and atmospheric distortion.
+ Other nights are included in the final data for analysis because while the data may not be entirely photometric. it remains reliable enough for color analysis because of the non-wavelength dependent nature of most remaining distortions.," Other nights are included in the final data for analysis because while the data may not be entirely photometric, it remains reliable enough for color analysis because of the non-wavelength dependent nature of most remaining distortions."
+ We discuss the errors in our derived SDSS colors further in refsec:colors, We discuss the errors in our derived SDSS colors further in \\ref{sec:colors1}.
+ Many of ]..the targets are luminous enough in red wavelengths that longer exposure times cause parts of the spectrum to saturate., Many of the targets are luminous enough in red wavelengths that longer exposure times cause parts of the spectrum to saturate.
+ The wavelengths from onward are of particular concern., The wavelengths from onward are of particular concern.
+ However. shorter exposure times limited the amount of information we could derive from wavelengths below because of low signal-to-noise ratios.," However, shorter exposure times limited the amount of information we could derive from wavelengths below because of low signal-to-noise ratios."
+ To create an S star spectral atlas. when possible we interpolated between two different exposure times: a shorter exposure time that produced accurate. unsaturated spectra at the longer wavelengths. and a longer exposure that resulted in higher signal-to-noise at shorter wavelengths.," To create an S star spectral atlas, when possible we interpolated between two different exposure times: a shorter exposure time that produced accurate, unsaturated spectra at the longer wavelengths, and a longer exposure that resulted in higher signal-to-noise at shorter wavelengths."
+ Generally. we determined by hand which spectra were reliable in a particular wavelength region by comparing multiple exposure times and examining long exposures for nonlinear behavior in areas from onward.," Generally, we determined by hand which spectra were reliable in a particular wavelength region by comparing multiple exposure times and examining long exposures for nonlinear behavior in areas from onward."
+ For stars with only one available exposure time. we rejected a star as saturated if the counts exceeded 60.000.," For stars with only one available exposure time, we rejected a star as saturated if the counts exceeded 60,000."
+ After interpolating between spectra to eliminate saturated regions and rejecting some data. we were left with 46 S star spectra.," After interpolating between spectra to eliminate saturated regions and rejecting some data, we were left with 46 S star spectra."
+ We use S star line identifications (Wyckoff&Clegg1978) to identify major spectral features and select objects with unique characteristics to include in the final spectral atlas. which includes 14 objects marked in reftab:stars..," We use S star line identifications \citep{1978MNRAS.184..127W} to identify major spectral features and select objects with unique characteristics to include in the final spectral atlas, which includes 14 objects marked in \\ref{tab:stars}."
+ The objects were chosen to be a representative sample of the larger FAST sample. which includes many fundamentally similar objects.," The objects were chosen to be a representative sample of the larger FAST sample, which includes many fundamentally similar objects."
+ Two sample spectra with associated molecular line identifications are plotted in reffig:speclines.., Two sample spectra with associated molecular line identifications are plotted in \\ref{fig:speclines}.
+ While the spectra display the same molecular bands. most notably those associated with ZrO. they have different overall colors. which may be due to differences in abundances and/or effective temperatures.," While the spectra display the same molecular bands, most notably those associated with ZrO, they have different overall colors, which may be due to differences in abundances and/or effective temperatures."
+ We also note that these spectra show some TiO bands. indicating that they are not pure S stars. but instead probably belong to the MS classification.," We also note that these spectra show some TiO bands, indicating that they are not pure S stars, but instead probably belong to the MS classification."
+ The appearance of TiO bands indicates the possibility of a C/O ratio slightly less than 0.95., The appearance of TiO bands indicates the possibility of a C/O ratio slightly less than 0.95.
+ The creation of a spectral atlas allows for greater investigation into the similarities and differences within the S star class., The creation of a spectral atlas allows for greater investigation into the similarities and differences within the S star class.
+" For instance. some members of the spectral atlas exhibit. in addition to prominent ZrO features. Ho. H7, and H7 emission lines."," For instance, some members of the spectral atlas exhibit, in addition to prominent ZrO features, $\alpha$, $\beta 
+$, and $\gamma$ emission lines."
+ The presence of these emission lines indicates that these objects are almost certainly AGB stars. e.g.. Miras.," The presence of these emission lines indicates that these objects are almost certainly AGB stars, e.g., Miras."
+ The spectral atlas is made available digitally as 14 individual FITS files via this journal., The spectral atlas is made available digitally as 14 individual FITS files via this journal.
+ We can immediately use the first digital spectral atlas of S stars to derive some medium-resolution spectral indices. potentially useful to classify stars as M. S. or C even where they may have very similar broadband colors.," We can immediately use the first digital spectral atlas of S stars to derive some medium-resolution spectral indices, potentially useful to classify stars as M, S, or C even where they may have very similar broadband colors."
+ Apart from our own FAST spectra of 48 S stars. we use M giants from the Indo-US spectral atlas (ValdesΕ..etal.2004).," Apart from our own FAST spectra of 48 S stars, we use M giants from the Indo-US spectral atlas \citep{Valdes04}."
+. We use eight carbon star spectra also obtained with the FAST for a separate project (Green et al., We use eight carbon star spectra also obtained with the FAST for a separate project (Green et al.
+ 2011. in preparation).," 2011, in preparation)."
+ These latter spectra span a wide range of spectral band strengths. but are selected from the SDSS color wedge defined by Margon.B..et(2002).," These latter spectra span a wide range of spectral band strengths, but are selected from the SDSS color wedge defined by \citet{Margon02}."
+ We defined three indices. one each for TiO. ZrO. and Cs.," We defined three indices, one each for TiO, ZrO, and $_2$."
+ We follow the basic premise of using ratios of mean flux per Ángstrómm across key molecular bands. so that spectra of differing resolution should yield similar results.," We follow the basic premise of using ratios of mean flux per ngströmm across key molecular bands, so that spectra of differing resolution should yield similar results."
+ We use a neighboring comparison region for each band. so that sensitivity to broadband flux calibration is minimized.," We use a neighboring comparison region for each band, so that sensitivity to broadband flux calibration is minimized."
+ For the ZrO index. we divide the mean flux across 6400 — (ZrO off-band) by that between 6475 — (ZrO band).," For the ZrO index, we divide the mean flux across 6400 – (ZrO off-band) by that between 6475 – (ZrO band)."
+ For the TiO index. we use mean Av! flux from 7065 — (TiO band) divided by 6965 — (TIO off-band).," For the TiO index, we use mean $^{-1}$ flux from 7065 – (TiO band) divided by 6965 – (TiO off-band)."
+ The C index is derived from 5025 — (Cs band) and 5238 — (off-band)., The C index is derived from 5025 – $_2$ band) and 5238 – (off-band).
+ Thus. lower flux in the denominator caused by strong molecular bands yields a larger spectral index value.," Thus, lower flux in the denominator caused by strong molecular bands yields a larger spectral index value."
+" We note that the C index near can also be substantially contaminated by TiO. so we label it ""CT517 hereafter."," We note that the C index near can also be substantially contaminated by TiO, so we label it “CT51” hereafter."
+ At least within this representative sample of stars. a combination of these indices yields à preliminary S star classification. as shown in Figure 2..," At least within this representative sample of stars, a combination of these indices yields a preliminary S star classification, as shown in Figure \ref{fig:specind}."
+ The objects most reliably classified as S stars would have ZrO index >1.2 and C» index <L.4., The objects most reliably classified as S stars would have ZrO index $>1.2$ and $_2$ index $<1.4$.
+ To use the SDSS database to find likely candidates for both S giants anddwarfs. we need reliable colors for our objects in as many SDSS bandpasses as possible.," To use the SDSS database to find likely candidates for both S giants anddwarfs, we need reliable colors for our objects in as many SDSS bandpasses as possible."
+ The spectra from FAST cover the bandpasses for the g and r filters completely. and we are thus able to convolve our spectra with these bandpasses and generate SDSS colors.," The spectra from FAST cover the bandpasses for the $g$ and $r$ filters completely, and we are thus able to convolve our spectra with these bandpasses and generate SDSS colors."
+ Each of the FAST spectra contains 2681 data points. with approximate wavelength separation of 1.47 A.," Each of the FAST spectra contains 2681 data points, with approximate wavelength separation of 1.47 ."
+. To convolve the spectra with the filter transmission curves. we use linear interpolation," To convolve the spectra with the filter transmission curves, we use linear interpolation"
+((Nova Per 1901: Campbell 1903). belongs to a subgroup of cataclysmic variables (CVs) called Intermediate: Polars (IPs).,"(Nova Per 1901; Campbell 1903), belongs to a subgroup of cataclysmic variables (CVs) called Intermediate Polars (IPs)."
+ In. these systems. an asvnchronously-rotating. magnetic white dwarl accretes material [rom a less-massive. ate-type companion filling its Roche lobe.," In these systems an asynchronously-rotating, magnetic white dwarf accretes material from a less-massive, late-type companion filling its Roche lobe."
+ Cas leaving the Companion star attempts to form an accretion disc around he primary star but its magnetic field either prevents the ormation of the cise or truncates it near the white chart., Gas leaving the companion star attempts to form an accretion disc around the primary star but its magnetic field either prevents the formation of the disc or truncates it near the white dwarf.
+ wawas identified with the X-ray source 0327|43 by line. ticketts Warwick (1979) and confirmed as an LP? by the detection of a 351ss X-ray spin pulse by Watson. Kine Osborne (1985: hereafter WIxXO) and Norton. Watson lxing (1988).," was identified with the X-ray source A0327+43 by King, Ricketts Warwick (1979) and confirmed as an IP by the detection of a s X-ray spin pulse by Watson, King Osborne (1985; hereafter WKO) and Norton, Watson King (1988)."
+ The same period was subsequently found. in optical photometry by. Patterson (1991)., The same period was subsequently found in optical photometry by Patterson (1991).
+" hhas the longest orbital period from the sample of known CVs. Po, = 2dd. (Crampton. Cowley Fisher 1986: hereafter CCE)."," has the longest orbital period from the sample of known CVs, $_{\mbox{orb}}$ = d, (Crampton, Cowley Fisher 1986; hereafter CCF)."
+ Phe wide binary separation combined with a relatively weak magnetic field (~ 1 MG) means that a truncated aceretion disc must be present if current theories of disc formation are correct (Llameury. ling Lasota 1986).," The wide binary separation combined with a relatively weak magnetic field $\sim$ 1 MG) means that a truncated accretion disc must be present if current theories of disc formation are correct (Hameury, King Lasota 1986)."
+ The presence of a disc has vet to be confirmed by direct. observation. although the system. does. undergo dwarf nova outbursts every 3vyvears. where its optical brightness increases from 13th to LOth magnitude (Sabbacdin Bianehini 1983).," The presence of a disc has yet to be confirmed by direct observation, although the system does undergo dwarf nova outbursts every years where its optical brightness increases from 13th to 10th magnitude (Sabbadin Bianchini 1983)."
+ The most-likely mechanism for cwart nova outbursts is a thermal instability within an accretion disc (Osaki 1974)., The most-likely mechanism for dwarf nova outbursts is a thermal instability within an accretion disc (Osaki 1974).
+ ooutbursts have been modelled: as such by Cannizzo Ixenvon (1986) and Kim. Wheeler Mineshige (1992).," outbursts have been modelled as such by Cannizzo Kenyon (1986) and Kim, Wheeler Mineshige (1992)."
+ This paper is à continuation of paper (Alorales-Rueda. sull Roche 1996). in which we presented spectrophotometric observations of ttaken on the rise to its 1996 outburst (Mattei 1996).," This paper is a continuation of paper (Morales-Rueda, Still Roche 1996), in which we presented spectrophotometric observations of taken on the rise to its 1996 outburst (Mattei 1996)."
+ We reported the detection of quasi-periodic oscillations (QPOs) within the Doppler-broadened emission lines of LL and Le11., We reported the detection of quasi-periodic oscillations (QPOs) within the Doppler-broadened emission lines of H and He.
+ Phis provides an opportunity to map the velocity structure of the oscillations., This provides an opportunity to map the velocity structure of the oscillations.
+ QPOs are defined as. Iow-coherence brightness oscillations thought to be associated with material within the inner accretion [lows of CVs., QPOs are defined as low-coherence brightness oscillations thought to be associated with material within the inner accretion flows of CVs.
+ Theoretical models developed to explain QPOs consider the presence of dense blobs of material orbiting in the inner regions of the accretion disc (Bath 1973). or non-racial pulsations over the surface of the white dwarf (Papaloizou," Theoretical models developed to explain QPOs consider the presence of dense blobs of material orbiting in the inner regions of the accretion disc (Bath 1973), or non-radial pulsations over the surface of the white dwarf (Papaloizou"
+spectra. one of them was summed (ice). so that the correct spectrum obtained from OBS! shows an excess with respect to the thermal component with the average gas temperature measured by (8.112:0.07. 9095: David 1993) at a somewhat lower confidence level of ~3.40 (see Table 1).,"spectra, one of them was summed twice), so that the correct spectrum obtained from OBS1 shows an excess with respect to the thermal component with the average gas temperature measured by $\pm$ 0.07,; David 1993) at a somewhat lower confidence level of $\sim 3.4\sigma$ (see Table 1)."
+ The fit with a single temperature gives ~9.91.3(71 keV. above je average gas temperature nieasured by (will a fiekl of view comparable to that of 1e PDS). implving the presence of a second spectral component.," The fit with a single temperature gives $\sim 9.9^{+1.3}_{-1.1}$ keV, above the average gas temperature measured by (with a field of view comparable to that of the PDS), implying the presence of a second spectral component."
+ The fit with two thermal 'oónponents (one fixed at 8.1 keV) requires an unrealistic second temperature (> 50 keV) wal stronglv supports a mechanism for the additional component. present in je spectrum of the Coma cluster., The fit with two thermal components (one fixed at 8.1 keV) requires an unrealistic second temperature $>$ 50 keV) that strongly supports a mechanism for the additional component present in the spectrum of the Coma cluster.
+ 1 we consider a power-law lor the second component. 1e PDS data are not able to fix the photon index. but the flux is rather stable against index variations.," If we consider a power-law for the second component, the PDS data are not able to fix the photon index, but the flux is rather stable against index variations."
+ We assume a photon index Dy. = 2.0 (o derive the fux iab results to be (2.341.0)x10|ergem7s+ in the 2080 keV οποιον range., We assume a photon index $\Gamma_X$ = 2.0 to derive the flux that results to be $(2.3\pm 1.0)\times 10^{-11}\erg$ in the 20–80 keV energy range.
+ The observed count rate of ODS2 is 0.7240.02 c(s/s in the 15100 keV energv range. at the confidence level of ~366.," The observed count rate of OBS2 is $\pm$ 0.02 cts/s in the 15–100 keV energy range, at the confidence level of $\sim
+36\sigma$."
+ Al energies above 20 keV the spectrum shows an excess with respect to the thermal emission (KT = 8.1 keV) at a confidence level of ~3.40. (see Table 1)., At energies above 20 keV the spectrum shows an excess with respect to the thermal emission (kT = 8.1 keV) at a confidence level of $\sim 3.4\sigma$ (see Table 1).
+ The fit with a single temperature gives ~9.5(5 keV. Also ODS2 indicates the presence of an additional spectral feature and also in this case the fit with a second thermal component requires unrealistic values for the temperature., The fit with a single temperature gives $\sim 9.5^{+0.8}_{-0.6}$ keV. Also OBS2 indicates the presence of an additional spectral feature and also in this case the fit with a second thermal component requires unrealistic values for the temperature.
+ The flux (Dy —2.0) is (13(5)x10.Hergem7?«!1 in the band 2080 keV. consistent with the flux reported in ODSI1.," The flux $\Gamma_X$ =2.0) is $(1.3^{+0.5}_{-0.6})\times 10^{-11}\erg$ in the band 20–80 keV, consistent with the flux reported in OBS1."
+ The fIuxes are (marginally) consistent also at a conlidence level: (2.30.7)x10I!ergem7s1 in the first observation and “ereem7s! in the second one., The fluxes are (marginally) consistent also at a confidence level: $(2.3\pm 0.7)\times 10^{-11}\erg$ in the first observation and $(1.3^{+0.3}_{-0.4})\times 10^{-11}\erg$ in the second one.
+ The combined spectrum is obtained by sunming (he spectra of the (wo observations (see Fie., The combined spectrum is obtained by summing the spectra of the two observations (see Fig.
+ 1)., 1).
+ The total count rate is 0.74040.017 cts/s in the 15100 keV energy range. αἱ the confidence level of ~44o.," The total count rate is $0.740\pm
+0.017$ cts/s in the 15–100 keV energy range, at the confidence level of $\sim 44\sigma$."
+ Al energies 220 keV the ΗΝ excess is at the confidence level of ~4.80 (see Table 1)., At energies $>$ 20 keV the HXR excess is at the confidence level of $\sim 4.8\sigma$ (see Table 1).
+ Even the inclusion of a svstematic to the data. necessary for sources with high S/N/ ratio but not [or [nint sources like Coma (Frontera 1991b). does not change the significance of our IINR excess.," Even the inclusion of a systematic to the data, necessary for sources with high S/N ratio but not for faint sources like Coma (Frontera 1997b), does not change the significance of our HXR excess."
+ The fit with a single thermal component gives 9.70.6 keV. well above the average gas temperature measured byGinga. wilh a statistically unacceptable 4? value (—2.1 for 8 d.o.f.).," The fit with a single thermal component gives $\pm0.6$ keV, well above the average gas temperature measured by, with a statistically unacceptable $\chi^2$ value (=2.1 for 8 d.o.f.)."
+ The presence of a second component is more evident from the 4? value that has a signilicant decrement when a second component. a power law. is added to the thermal component with KT—8.1 keV. The improvement passing from the first model (42 = 4.10 for 9 d.o.L)," The presence of a second component is more evident from the $\chi^2$ value that has a significant decrement when a second component, a power law, is added to the thermal component with kT=8.1 keV. The improvement passing from the first model $\chi^2_\nu$ = 4.10 for 9 d.o.f.)"
+ to the second one (47 = 1.2 for 7 d.o.L.), to the second one $\chi^2_\nu$ = 1.2 for 7 d.o.f.)
+ is significant at more than confidence level. according to the F-test.," is significant at more than confidence level, according to the F-test."
+ Also. the combined spectrum cannot be fitted with a second thermal component unless an unrealistic value for the temperature is assunied. (hus supporting the origin for," Also, the combined spectrum cannot be fitted with a second thermal component unless an unrealistic value for the temperature is assumed, thus supporting the origin for"
+(or even more elficient) in exchanging enerev with the equilibrium flow.,(or even more efficient) in exchanging energy with the equilibrium flow.
+ We know from previous studies that in the absence of the gravitv-induced stratification among the MITD waves. the fast magnetosonic waves are (he most efficient ones in exchanging enerev with (he equilibrium flow (Poedts.Rogava&Mahajanivan 2000).," We know from previous studies that in the absence of the gravity-induced stratification among the shear-modified MHD waves, the fast magnetosonic waves are the most efficient ones in exchanging energy with the equilibrium flow \citep{prm99,rpm00}."
+. It is of great. interest to check whether the fast Gravito-magnetosonic waves vive (he same quality. and (o determine what their contribution in the angular momentum redistribution could be., It is of great interest to check whether the fast Gravito-magnetosonic waves have the same quality and to determine what their contribution in the angular momentum redistribution could be.
+ The study of the linear. dynamics of all Gravito-MIID. waves is currently initiated and the results will be published in a subsequent paper., The study of the linear dynamics of all Gravito-MHD waves is currently initiated and the results will be published in a subsequent paper.
+ Finally. one has to remember (hat the non-modal approach that has been applied in the present paper. provides no information about the spatial aspects of the shear-induced processes. because (he study of (he Spatial Fourier LLarmonics is confined to the phase space of the wave number vectors k(/).," Finally, one has to remember that the non-modal approach that has been applied in the present paper, provides no information about the spatial aspects of the shear-induced processes, because the study of the Spatial Fourier Harmonics is confined to the phase space of the wave number vectors ${\bf k}(t)$."
+ In order to have a clear idea about the spatial appearance of (hese processes. one has to study (hem numerically (similarly to the uniform. non-stratilied flow study made by Bodoetal... (2001))) and to check how the non-mocdal phenomena couple with the traditional rotational instabilities.," In order to have a clear idea about the spatial appearance of these processes, one has to study them numerically (similarly to the uniform, non-stratified flow study made by \citet{b01}) ) and to check how the non-modal phenomena couple with the traditional rotational instabilities."
+ These results were obtained in the framework of the projects GOA 2004/01 (IX.U.Leuven). G.0304.07 (FWO-Vlaanderen) aud C90203 (ESA Prodex 8).," These results were obtained in the framework of the projects GOA 2004/01 (K.U.Leuven), G.0304.07 (FWO-Vlaanderen) and C90203 (ESA Prodex 8)."
+ Andria Rogava wishes to thank the (Leuven. Belgium) and the Physics(CIrieste. Italv) for supporting him. in part. through a Senior Postdoctoral Fellowship and Senior Associate Membership Award. respectively.," Andria Rogava wishes to thank the (Leuven, Belgium) and the (Trieste, Italy) for supporting him, in part, through a Senior Postdoctoral Fellowship and Senior Associate Membership Award, respectively."
+ The research ol Andria Rogava and Grigol Gogoberidze was supported in part by the Georgian National science Foundation grant GNSE/ST06/4-096., The research of Andria Rogava and Grigol Gogoberidze was supported in part by the Georgian National Science Foundation grant GNSF/ST06/4-096.
+ The research of Grigol Gogoberidze was supported in part bv the INTAS erant. 06-10000177-9258., The research of Grigol Gogoberidze was supported in part by the INTAS grant 06-1000017-9258.
+High-mass stars (OB spectral type. M>8M. and L10° Ls). although few in number. play a major role in the energy budget of galaxies. through their radiation. wind and the supernovae.,"High-mass stars (OB spectral type, $M>8\,M_{\odot}$ and $L>10^3\,L_{\odot}$ ), although few in number, play a major role in the energy budget of galaxies, through their radiation, wind and the supernovae."
+ They are believed to form by accretion in dense cores within molecular cloud complexes2002::2003: 2005) and/or coalescence (e.g. 2001))., They are believed to form by accretion in dense cores within molecular cloud complexes; ) and/or coalescence (e.g. ).
+ The intense radiation field emitted by a newly-formed central star heats and ioizes its parental molecular cloud. leading to the formation of a hot core (HC. e.g. 1997)) and afterwards an HII region.," The intense radiation field emitted by a newly-formed central star heats and ionizes its parental molecular cloud, leading to the formation of a hot core (HC, e.g. ) and afterwards an HII region."
+ Our current understanding of their formation remains poor. especially concerning the earliest phases of the process.," Our current understanding of their formation remains poor, especially concerning the earliest phases of the process."
+ The main observational difficulty is that high-mass stars are fewer in number thar low-mass stars and the molecular clouds that are able to form high-mass stars are statistically more distant than those forming low-mass stars., The main observational difficulty is that high-mass stars are fewer in number than low-mass stars and the molecular clouds that are able to form high-mass stars are statistically more distant than those forming low-mass stars.
+ Therefore. current observational studies of high-mass star formatiot suffer both from the lack of spatial resolution and. consequently. from a lack of theoretical understanding.," Therefore, current observational studies of high-mass star formation suffer both from the lack of spatial resolution and, consequently, from a lack of theoretical understanding."
+ The high-mass star-forming region G19.61-0.243 is an interesting target for the study of star cluster formation given its richness in terms of young stellar objects (YSOs). as indicated by studies at centimeter (e.g.1998:: 2000)). millimeter (e.g. 2005)) and (MIR: 2003)) wavelengths.," The high-mass star-forming region G19.61-0.23 is an interesting target for the study of star cluster formation given its richness in terms of young stellar objects (YSOs), as indicated by studies at centimeter (e.g.; ), millimeter (e.g. ) and mid-infrared (MIR; ) wavelengths."
+ is located at a distance of 12.6 kpe (see 2003)). based on the 21 em HI absorption spectrum toward the source.," G19.61-0.23 is located at a distance of 12.6 kpc (see ), based on the 21 cm HI absorption spectrum toward the source."
+ The total bolometric luminosity 1997)) is about 2x10°L...," The total bolometric luminosity ) is about $2\times10^6\,L_{\odot}$."
+" The region contains OH1983:; 1989)), water 1996)) and methanol1995:;2000:: 2000)) masers. and a grouping of UC HII regions and extended radio continuum. indicating that it is an active region of massive star formation."," The region contains OH; ), water ) and methanol; ) masers, and a grouping of UC HII regions and extended radio continuum, indicating that it is an active region of massive star formation."
+ The radio continuum emission from this region comes from five main sources. all of which are discussed in detail in (1998).," The radio continuum emission from this region comes from five main sources, all of which are discussed in detail in ."
+. The region has been mapped in CS. NH; and CO1999::1992:: 1998)).," The region has been mapped in CS, $_3$ and CO; )."
+ Extended mid- emission associated with the region is also detected by the MSX satellite 2003))., Extended mid-infrared emission associated with the region is also detected by the MSX satellite ).
+ Single-dish observations of molecular lines with high critical densities show the presence of dense molecular gas over a broad range in velocity 1992))., Single-dish observations of molecular lines with high critical densities show the presence of dense molecular gas over a broad range in velocity ).
+ In this paper we present a spectroscopic study of the region surrounding G19.61-0.23 in several transitions of CO isotopologues at an angular resolution of ~46” (about 2.8 pc at the distance of 12.6 kpc). on a large region of about 23x23’ (roughly 85 pe) centered on G19.61-0.23.," In this paper we present a spectroscopic study of the region surrounding G19.61-0.23 in several transitions of CO isotopologues at an angular resolution of $\sim46^{\prime\prime}$ (about 2.8 pc at the distance of 12.6 kpc), on a large region of about $23^{\prime}\times23^{\prime}$ (roughly 85 pc) centered on G19.61-0.23."
+ Millimeter observations of carbon monoxide provide useful information on the physical properties of dense interstellar clouds as well as on their dynamical state., Millimeter observations of carbon monoxide provide useful information on the physical properties of dense interstellar clouds as well as on their dynamical state.
+ Moreover. supplementary measurements of continuum emission in the sub-millimeter range with APEX. based on ATLASGAL data. and in the mid-infrared with Spitzer. based on GLIMPSE and MIPSGAL data. are presented.," Moreover, supplementary measurements of continuum emission in the sub-millimeter range with APEX, based on ATLASGAL data, and in the mid-infrared with Spitzer, based on GLIMPSE and MIPSGAL data, are presented."
+ The aim is to study the global scale physical properties and their relation with the small-scale characteristics., The aim is to study the global large-scale physical properties and their relation with the small-scale characteristics.
+ We analyzed the physical conditions and the velocity structure of the molecular components across the region., We analyzed the physical conditions and the velocity structure of the molecular components across the region.
+ We finally consider the possible implications for, We finally consider the possible implications for
+"In summary. (he ordering of Gime scales (hat will affect the physics of supernovae is roughly: where Tyl ms is the dynamical time scale. 7,,, is the (sub-Ixeplerian) rotation (niescale. 744,por ο rotation times. is (he timescale for dynamical bar formation. ms is the time lor the shock to form and stall. 74554~30 ms is the time for the MRI to erow the magnetic field to saturation. Tyεν~50—100 ms is the time for non-axisvmmetrie an = 1) modes (to grow. Τρ»0.1— Ls is the time to spin down due to magnetosonic huninosity. Τωρα0.1— Ls is the time lor a secular bar mode to grow. 7;4,,~1—10s is the de-leptonization time of the PNS to contract to [orm a neutron stab. 77,4655,~LO 8 is (he time lor the successful shock to propagate oul of the infalling iron core ancl into the surrounding star. e. g.. the helium core. and τος, 15 the time to dissipate angular momentum and rotational energv of the core by gravitational radiation reaction forces.","In summary, the ordering of time scales that will affect the physics of supernovae is roughly: where $\tau_{dyn} \sim 1 $ ms is the dynamical time scale, $\tau_{rot}$ is the (sub-Keplerian) rotation timescale, $\tau_{dyn-bar}$, a few rotation times, is the timescale for dynamical bar formation, $\tau_{shock} \sim 10$ ms is the time for the shock to form and stall, $\tau_{MRI} \sim 30 $ ms is the time for the MRI to grow the magnetic field to saturation, $\tau_{NAXI} \sim 50 -100$ ms is the time for non-axisymmetric (m = 1) modes to grow, $\tau_{mhd} \sim 0.1 - 1$ s is the time to spin down due to magnetosonic luminosity, $\tau_{secular} \sim 0.1 - 1$ s is the time for a secular bar mode to grow, $\tau_{delep} \sim 1 - 10 s$ is the de-leptonization time of the PNS to contract to form a neutron star, $\tau_{explosion} \sim 10$ s is the time for the successful shock to propagate out of the infalling iron core and into the surrounding star, e. g., the helium core, and $\tau_{grr}$ is the time to dissipate angular momentum and rotational energy of the core by gravitational radiation reaction forces."
+ In the current context. Trrpfosion 15 meant to be the Gime bevond which physical processes in the PNS will no longer affect the ultimate outcome. perhaps because the density has decreased sufficiently that even if magnetoacoustic Πας is liberated. it cannot propagate outward aud hence affect the explosion.," In the current context, $\tau_{explosion}$ is meant to be the time beyond which physical processes in the PNS will no longer affect the ultimate outcome, perhaps because the density has decreased sufficiently that even if magnetoacoustic flux is liberated, it cannot propagate outward and hence affect the explosion."
+ In the absence of a detailed model of this process. we have taken (he time for a successful shock to propagate into the helium core as a representative measure of this scale.," In the absence of a detailed model of this process, we have taken the time for a successful shock to propagate into the helium core as a representative measure of this scale."
+ We note that the epoch of the onset of convection within the PNS is of order of tens of ms and that of convection in the post-shock region is of order 100 ms., We note that the epoch of the onset of convection within the PNS is of order of tens of ms and that of convection in the post-shock region is of order 100 ms.
+ This emphasizes that (imescales of order 0.1 to 1 s. not epochs of tradiüonal concentration. will involve a varletv. of interacting physical processes (hat. will need to be more deeply understood.," This emphasizes that timescales of order 0.1 to 1 s, not epochs of traditional concentration, will involve a variety of interacting physical processes that will need to be more deeply understood."
+" We also note that the timescale τος, is long. so all the physics associated with the other time scales represented here must be solved to know the conditions that might be relevant to the production of gravity waves once (if ever) those become (he dominant sink."," We also note that the timescale $\tau_{grr}$ is long, so all the physics associated with the other time scales represented here must be solved to know the conditions that might be relevant to the production of gravity waves once (if ever) those become the dominant sink."
+ The contraction phase (Burrows&Lattimer1956:NeilJanka1995;Villainetal.2004) will lead to spin up and perhaps to crossing the threshold for NANI or enhancing the growth rate of these instabilities. (he amplitude of which max depend on T/Wl.," The contraction phase \citep{bur86,kei95,pon99,vil04} will lead to spin up and perhaps to crossing the threshold for NAXI or enhancing the growth rate of these instabilities, the amplitude of which may depend on $\tw$."
+ These instabilities will cause some loss of rotation energy and angular momentum as (he core contracts. perhaps altering the specific nonaxisvuunetric modes (hat come into play.," These instabilities will cause some loss of rotation energy and angular momentum as the core contracts, perhaps altering the specific non–axisymmetric modes that come into play."
+ The core will dissipate its differential rotation and angular momentum until the loss rales become comparable to. or longer than. the contraction (time scale.," The core will dissipate its differential rotation and angular momentum until the loss rates become comparable to, or longer than, the contraction time scale."
+ The celeptonization, The de–leptonization
+a sinele average before performing aperture plotomietry ou UY Vol aud a comparison star (both shown in Fie. 1)).,a single average before performing aperture photometry on UY Vol and a comparison star (both shown in Fig. \ref{ImageFig}) ).
+ All photometry was done differentially relative to the comparison star using standard tecbuiques. and this star was then calibrated separately relative to standard stars A.C. aud D in the field of T Phe observed ou inultiple photometric nights (Landolt1992)..," All photometry was done differentially relative to the comparison star using standard techniques, and this star was then calibrated separately relative to standard stars A, C, and D in the field of T Phe observed on multiple photometric nights \citep{Landolt:1992a}."
+. Our estimate of the calibrated magnitude of the comparison is R=11.66-20.03., Our estimate of the calibrated magnitude of the comparison is $R=14.66\pm0.03$.
+ The uncertainty quoted is the might-to-might standard deviation., The uncertainty quoted is the night-to-night standard deviation.
+ Systematic errors may be larger as color corrections were not possible since only &R baud observatious were performecl., Systematic errors may be larger as color corrections were not possible since only $R$ band observations were performed.
+ Om deduced average magnitude for UY Vol is R=22.39c 0.01., Our deduced average magnitude for UY Vol is $R=22.39\pm0.04$ .
+ This is the formal error on the mean and systematic uncertaiutfies in the calibration may be larger., This is the formal error on the mean and systematic uncertainties in the calibration may be larger.
+ We show the long-term Lehtcurve iu Fie. 2.., We show the long-term lightcurve in Fig. \ref{LongLCFig}.
+ Considerable nieht-to-uight variability is preseut but there is no obvious lone-terim trend., Considerable night-to-night variability is present but there is no obvious long-term trend.
+ We will quantity this statement after removing some of the intrinsic variability in Section [.., We will quantify this statement after removing some of the intrinsic variability in Section \ref{LightcurveSection}.
+ To verify the significance of the variability seen we also show a lightcurve for another nearby star at R=22.124001., To verify the significance of the variability seen we also show a lightcurve for another nearby star at $R=22.42\pm0.04$.
+" The standard deviation of the individual UY Vol data is πας, while that for the nuon-variable star is uma."," The standard deviation of the individual UY Vol data is mag, while that for the non-variable star is mag."
+ IR data were obtained iu a T-poiut dither pattern., IR data were obtained in a 7-point dither pattern.
+ ον 50ss images were taken at 6 of the 7 positious. and 2 at the other.," $3\times50$ s images were taken at 6 of the 7 positions, and 2 at the other."
+ For cach might. à sky image derived from the median of the dithered inages was subtracted. and flat-fields were applied.," For each night, a sky image derived from the median of the dithered images was subtracted, and flat-fields were applied."
+ We excluded: images with the highest sky values and those with a sky value sieuificautlv deviating from the nightly mean to minimize residuals in the background subtraction., We excluded images with the highest sky values and those with a sky value significantly deviating from the nightly mean to minimize residuals in the background subtraction.
+ We then filtered the remaining nuages based ou visibility of the faintest stars in the field., We then filtered the remaining images based on visibility of the faintest stars in the field.
+ Our final combination of these best nuages used 122 individual frames. all of which had beeu oeidividuallv checked.," Our final combination of these best images used 422 individual frames, all of which had been individually checked."
+ The target is mareinally detected oei this combined IR image at a position consistent with hat measured frou optical inages., The target is marginally detected in this combined IR image at a position consistent with that measured from optical images.
+ Photometry relative ο several 2ATASS stars iu the field vields 7=21.84 L2., Photometry relative to several 2MASS stars in the field yields $J=21.3\pm0.2$ .
+ At this level. we caution that svstematic errors in vackeround subtraction are likely to be larger than the cornu statistical error quoted.," At this level, we caution that systematic errors in background subtraction are likely to be larger than the formal statistical error quoted."
+ It was innucdiately apparent that the data appeared consistent with modulation ou the published orbital veriod of ανν = 0.1593«dd (Wolffetal.2002) with a snele-huniped modulation., It was immediately apparent that the data appeared consistent with modulation on the published orbital period of hrs = d \citep{Wolff:2002a} with a single-humped modulation.
+ To verity this we fitted the full dataset with a sinusoidal modulation of variable period. allowing the plasing. amplituce aud uean brightuess to vary freely.," To verify this we fitted the full dataset with a sinusoidal modulation of variable period, allowing the phasing, amplitude and mean brightness to vary freely."
+ We find several strong uinina iu the dd period ranee (Fie. 39)., We find several strong minima in the d period range (Fig. \ref{SineFitFig}) ).
+ One of these is consistent with the orbital period. aud the others are consistent with one-day aliases of the orbital xriod. as expected given our ouce-per-day sampling.," One of these is consistent with the orbital period, and the others are consistent with one-day aliases of the orbital period, as expected given our once-per-day sampling."
+ No sienificant minimuia are seen other than these aliases., No significant minima are seen other than these aliases.
+" Choosing the alias corresponding to the X-ray period. we derive an optical period of P=60.159351+000012 dd. The 1...@ uncertainty quoted corresponds o the range of periods with which 4?x:4Z,,|1."," Choosing the alias corresponding to the X-ray period, we derive an optical period of $P=0.159331\pm0.000012$ d. The $1-\sigma$ uncertainty quoted corresponds to the range of periods with which $\chi^2 \leq \chi^2_{min}+1$."
+ This 2ο1οςἱ is conusisteut with the secure A-rav orbital period of dd to within errors., This period is consistent with the secure X-ray orbital period of d to within errors.
+ The uucertaiutv is ~0.01 xeriod is mdeed orbital not sisuificautlv longer as would ο expected from a superluup modulation., The uncertainty is $\sim0.01$ period is indeed orbital not significantly longer as would be expected from a superhump modulation.
+"lis fainter than the normal stars with the same effective temperature in the far-UV on wavelengths lower than 2400A, whereas it is brighter than normal stars in the near-UV and visible regions.","is fainter than the normal stars with the same effective temperature in the far-UV on wavelengths lower than $2400\,$, whereas it is brighter than normal stars in the near-UV and visible regions."
+" Although the comparison of the narrow-band variations revealed the regions where the disagreement between the observed and predicted flux variations occurs, the narrow-band variations are inadequate to figure out the origin of these"," Although the comparison of the narrow-band variations revealed the regions where the disagreement between the observed and predicted flux variations occurs, the narrow-band variations are inadequate to figure out the origin of these"
+in this fit.,in this fit.
+ The recoil velocities as given in Eq. (1)), The recoil velocities as given in Eq. \ref{eqn:Fit3}) )
+ are plotted in Fig., are plotted in Fig.
+ b. as a functiou of mass ratio ancl spiLI parameter., \ref{Fig-v-recoil} as a function of mass ratio and spin parameter.
+ Iu order to compute the probability that au. IMDII ronuünus in its elobular cluster. our caleulation shall proceed as follows.," In order to compute the probability that an IMBH remains in its globular cluster, our calculation shall proceed as follows."
+ First we begim bv assunimg that an IMDITI has formed within a elobular cluster wit1 a particularinitial mass. A_yppiy. which we shall vary.," First we begin by assuming that an IMBH has formed within a globular cluster with a particularinitial mass, $M_{\mathrm{IMBH}}$, which we shall vary."
+ Second. we further assunie a certain mass distribution for the DIIs in the vicinity of the ΤΑΠΟΠ. which we shall also vary.," Second, we further assume a certain mass distribution for the BHs in the vicinity of the IMBH, which we shall also vary."
+ Third. we subject this IMIBIT to a number of mergers expected within a proto globular cluster environnent that we shall describe below.," Third, we subject this IMBH to a number of mergers expected within a proto globular cluster environment that we shall describe below."
+ Finally. we deteriuue the probability that the kick velocity for the IMDII has remained below the canonical globular cluster escape velocity (50kms 7) during the entire chain of nergers.," Finally, we determine the probability that the kick velocity for the IMBH has remained below the canonical globular cluster escape velocity $50 \KMS$ ) during the entire chain of mergers."
+ Even in the absence of an INDIT. DITs eject themselves yon globular clusters via standard few-hody interactions on a timescale of ~1Cyr after the onset of dass segregation (22?2)..," Even in the absence of an IMBH, BHs eject themselves from globular clusters via standard few-body interactions on a timescale of $\sim 1~{\rm Gyr}$ after the onset of mass segregation \citep{Kulkarni:93bhgc,Sigurdsson:93bhgc,Portegies:00bhmerge,Oleary:2005bm}."
+ Therefore. due to such Newtoniau ew-body interactions. the supply of DIIS is eveutually depleted.," Therefore, due to such Newtonian few-body interactions, the supply of BHs is eventually depleted."
+ With an IMDBIL however. this process speeds up lupressively. as ejections by interactions with an IMDII ecole the doninaut source of stellar-imass DII ejections (?)..," With an IMBH, however, this process speeds up impressively, as ejections by interactions with an IMBH become the dominant source of stellar-mass BH ejections \citep{Gultekin:2006tb}."
+ As iu most few-body interactions. the ejection of one object tighteus the orbit of a remaining bound pair. in this case an IMDBIL-DII binary and after several subsequent ejections. the INIBU-BU binary mcrecs.," As in most few-body interactions, the ejection of one object tightens the orbit of a remaining bound pair, in this case an IMBH-BH binary – and after several subsequent ejections, the IMBH-BH binary merges."
+ Soon iter all the DII« have been evacuated. the short epoch of IAIBU-BIT inerecrs euds.," Soon after all the BHs have been evacuated, the short epoch of IMBH-BH mergers ends."
+ Within this theoretical framework. it is possible to construct a fiducial nuniber of mergers for a proto elobular cluster.," Within this theoretical framework, it is possible to construct a fiducial number of mergers for a proto globular cluster."
+ This πανο can be written as Gultekin et al., This number can be written as Gultekin et al.
+ 2006 predict Mae~25 per IMDIT. and we adopt this for the fiducial number of mergers that the IMIBIT eucouuters.," 2006 predict $N_{\mathrm{merge}} \sim 25$ per IMBH, and we adopt this for the fiducial number of mergers that the IMBH encounters."
+ Although we do vary this parameter in figure 2.. the dependence of the retention probability on the wmmber of niergers is relatively iinor. since the IMIBIT erows in mass over cach merger aud the kick velocity increases with increasing mass ratio.," Although we do vary this parameter in figure \ref{Fig-ret-prob}, the dependence of the retention probability on the number of mergers is relatively minor, since the IMBH grows in mass over each merger and the kick velocity increases with increasing mass ratio."
+ Tn order to assign a kick velocity to each of tle merecrs. we choose the orieutation. spin. nass. and eccentricity.," In order to assign a kick velocity to each of the mergers, we choose the orientation, spin, mass, and eccentricity."
+ We outline the assumptions made for cach distribution below., We outline the assumptions made for each distribution below.
+ Let us first discuss the issue of the initial spin orientation., Let us first discuss the issue of the initial spin orientation.
+ Uvdrodvuamic interactions between a gas disk aud a black hole binary are believed to align the spin directions to the angular momentum axis of the binary orbital plane im many active galaxies (?).., Hydrodynamic interactions between a gas disk and a black hole binary are believed to align the spin directions to the angular momentum axis of the binary orbital plane in many active galaxies \citep{tamara:07spin}.
+ Towever. the euvironnieut of a globular cluster is not particularly eas-rich. so there is nomitio roason to expect the black role spins to be aligued.," However, the environment of a globular cluster is not particularly gas-rich, so there is no reason to expect the black hole spins to be aligned."
+ We therefore assuue au isotropic distribution of orientation angeles for cach encouuter., We therefore assume an isotropic distribution of orientation angles for each encounter.
+ Let us now discuss the choice of spin magnitude., Let us now discuss the choice of spin magnitude.
+ Most heories predict a non-zero spin for a black hole produced via stellar runaway (7) or frou a supernovae renuit (2).., Most theories predict a non-zero spin for a black hole produced via stellar runaway \citep{Rees:07mbh} or from a supernovae remnant \citep{Fryer:01spin}.
+ If au IMDBIT started with zero spin. a imerecr is likely o spin up the remnant through transfer of orbital to spin augular iiomenutuim (?)..," If an IMBH started with zero spin, a merger is likely to spin up the remnant through transfer of orbital to spin angular momentum \citep{Gammie:2003qi}."
+ Towever. a I&err black hole can spindowo when magnetic field lines thread through he ergosphere to magueticalle brake the svstein (?).. oovided there is a gaseous disk around the reumaut.," However, a Kerr black hole can $spin
+ down$ when magnetic field lines thread through the ergosphere to magnetically brake the system \citep{Blandford:77spin}, provided there is a gaseous disk around the remnant."
+ Taking all these consideration iuto account. we shall explore three cases: (1) the spin maguitude is selected roni a uniformi initial spin distribution (fiducial case): (2) the initial spin maeuitude of the IMDIT κους is set ο U.998MRrays and (3) the spin is initially set to zero.," Taking all these consideration into account, we shall explore three cases: (1) the spin magnitude is selected from a uniform initial spin distribution (fiducial case); (2) the initial spin magnitude of the IMBH seed is set to $0.998 M_{\mathrm{IMBH}}^2$; and (3) the spin is initially set to zero."
+ We asstune the spin of the secondary. DII to be randomly selected from a distribution of [0.0998]AL? where M. is the mass of the secoudary BID.," We assume the spin of the secondary BH to be randomly selected from a distribution of $[0,0.998]~M_{\mathrm{sec}}^2$ , where $M_{\mathrm{sec}}$ is the mass of the secondary BH."
+ Since these stelbu- DIIS originate as a supernova renimant. though. the," Since these stellar-mass BHs originate as a supernova remnant, though, the"
+Efforts in understanding the physics at work in active ealactic nuclei (AGN) started four decades ago.,Efforts in understanding the physics at work in active galactic nuclei (AGN) started four decades ago.
+ The origin ol the infrared (Hi) continuum of AGN was initially a matter of controversy. as it could be non-thermal but. could equally be due to thermal emission from dust. grains.," The origin of the infrared (IR) continuum of AGN was initially a matter of controversy, as it could be non-thermal but could equally be due to thermal emission from dust grains."
+ Η was long ago sugeested(Reesetal.1969). that Ht emission radiation from Sevlort galaxies in the 2.2. 22 wavelength range was produced by dust grains heated by ultraviolet (UV) and optical emission [from the nucleus.," It was long ago suggested\citep{rees69}
+ that IR emission radiation from Seyfert galaxies in the 2.2 – 22 wavelength range was produced by dust grains heated by ultraviolet (UV) and optical emission from the nucleus."
+ Work carried out. later on sugeests the LIU emission to be the reprocessed emission of the UV/optical radiation from the accretion disk by the particles composing the torus. namely silicate and graphite erains (e.g. Pier&Ixrolik.1992: Granato&Danese 1904: Efstathiou&Rowan-Robinson 1995: Nenkovaetal. 2002)).," Work carried out later on suggests the IR emission to be the reprocessed emission of the UV/optical radiation from the accretion disk by the particles composing the torus, namely silicate and graphite grains (e.g. \citealt{pier92}; \citealt{granato94}; \citealt{efstathiou95}; \citealt{nenkova02}) )."
+ Various configurations of the dust distribution gcometry and compositions havebeen since suggested (e.g. Pier&Ixrolik 1992: vanBenumel&Dullemond 2003:: Dullemond&van 2005::Fritzetal. 2006:: Elitzur&Shlosman 2006))., Various configurations of the dust distribution geometry and compositions havebeen since suggested (e.g. \citealt {pier92}; ; \citealt{vanbemmel03}; ; \citealt{dullemond05}; \citealt{fritz06}; ; \citealt{elitzur06}) ).
+ Recentobservations (Jalleetal.2004). indicate that, Recentobservations \citep{jaffe04} indicate that
+ollects of arbitrarily strong electromagnetic fields using a QED one-loop ellective Lagrangian approach.,effects of arbitrarily strong electromagnetic fields using a QED one-loop effective Lagrangian approach.
+ These effects are discussed. in section 2.2.., These effects are discussed in section \ref{sec:tensors}.
+ Plasma elfects. are included bv assuming free electrons moving under the Lorentz force without any selt-interactions., Plasma effects are included by assuming free electrons moving under the Lorentz force without any self-interactions.
+ The model is that. of a cold magnetohyelrodyvnamic plasma ancl is. discussed. in section 2.3..., The model is that of a cold magnetohydrodynamic plasma and is discussed in section \ref{sec:plasma}.
+ We have also assumed. that the medium is homogeneous in agreement with the travelling-wave ansatz., We have also assumed that the medium is homogeneous in agreement with the travelling-wave ansatz.
+ OL course the actual situation is more complicated: with a thermally excited. plasma (e.g.7) and. inhomogeneities the latter can result in a whole slew of interesting interactions between the wave modes (2???) that are especially crucial to our understanding. of the thermal radiation from their surfaces. but these are bevond the scope of this paper.," Of course the actual situation is more complicated with a thermally excited plasma \citep[e.g.][]{Gill09BW}
+ and inhomogeneities — the latter can result in a whole slew of interesting interactions between the wave modes \citep{Heyl99polar,Heyl01qed,Heyl01polar,2003PhRvL..91g1101L} that are especially crucial to our understanding of the thermal radiation from their surfaces, but these are beyond the scope of this paper."
+ The formation of electromagnetic shocks is expected to be an important phenomenon for electromagnetic waves in the magnetized vacuum since. electromagnetic: waves can evolve discontinuities under the inlluence of nonlinear interactions. (272)..," The formation of electromagnetic shocks is expected to be an important phenomenon for electromagnetic waves in the magnetized vacuum since electromagnetic waves can evolve discontinuities under the influence of nonlinear interactions \citep{PhysRev.113.1649,
+ zheleznyakov1982shock, heyl1998electromagnetic}."
+ Such. shocks can form even in the presence of a plasma (?7).., Such shocks can form even in the presence of a plasma \citep{PhysRevD.59.045005}.
+ 1n this study. through our explicit focus on travelling waves. we examine an alternate class of solutions to the wave equations which do not suller this fate.," In this study, through our explicit focus on travelling waves, we examine an alternate class of solutions to the wave equations which do not suffer this fate."
+" αφίσας, they are stabilized against the formation of discontinuities by nonlinear features."," Instead, they are stabilized against the formation of discontinuities by nonlinear features."
+ “Phese waves travel as periodic wave trains without anv change to their form. such as wave steepening or shock formation.," These waves travel as periodic wave trains without any change to their form, such as wave steepening or shock formation."
+ Waves such as these may contribute to the formation of pulsar microstructures (??)," Waves such as these may contribute to the formation of pulsar microstructures \citep{1987STIN...8816622C,2001ApJ...558..302J}."
+ The vacuum of QED in the presence of large. magnetic fields can be described as a non-linear optical medium (?).., The vacuum of QED in the presence of large magnetic fields can be described as a non-linear optical medium \citep{1997JPhA...30.6485H}.
+ We also choose to treat the ellect of the plasma on the waves through source terms py and Ji: therefore. we begin bv considering. Alaxwell’s equations in the presence. of a medium and plasma sources.," We also choose to treat the effect of the plasma on the waves through source terms $\rho_p$ and $\Jp$; therefore, we begin by considering Maxwell's equations in the presence of a medium and plasma sources."
+ In. Lleavisicle-Lorentz units with e=1. Maxwell's equations can be used to derive the wave equations For clarity. we will avoid making cancellations or dropping vanishing terms.," In Heaviside-Lorentz units with $c=1$, Maxwell's equations can be used to derive the wave equations For clarity, we will avoid making cancellations or dropping vanishing terms."
+ We cleline the vacuum cliclectric and inverse magnetic permeability tensors as follows (?) In the next few sections we build a mocdel describing travelling waves in a magnetar's atmosphere from. these equatIons.," We define the vacuum dielectric and inverse magnetic permeability tensors as follows \citep{Jack75}
+ In the next few sections we build a model describing travelling waves in a magnetar's atmosphere from these equations."
+ In this section. we describe our model of the QED vacuum in strong background. fields in terms of vacuum cielectric and inverse magnetic permeability tensors.," In this section, we describe our model of the QED vacuum in strong background fields in terms of vacuum dielectric and inverse magnetic permeability tensors."
+ These are most conveniently described in terms of two Lorentz invariant combinations of the fields., These are most conveniently described in terms of two Lorentz invariant combinations of the fields.
+ In order to examine the nonlinear ellects of the vacuum. nonperturbatively. we wish to use vacuum ciclectric aud inverse magnetic permeability tensors which are valid to all orders in the fields.," In order to examine the nonlinear effects of the vacuum nonperturbatively, we wish to use vacuum dielectric and inverse magnetic permeability tensors which are valid to all orders in the fields."
+ Analytic expressionsfor these tensors were derived by? for the case of wrenchless fields (A.=(4E.BY?= 0) from the LUcisenbere-Euler-Weisskopl- (??7) one-loop ellective Lagrangian in ? and expressed in terms of a set of analytic functions.," Analytic expressionsfor these tensors were derived by \citet{1997JPhA...30.6485H} for the case of wrenchless fields $K=-(4\efi \cdot \bfi)^2=0$ ) from the Heisenberg-Euler-Weisskopf-Schwinger \citep{heisenberg-1936-98,
+ weisskopf1936kongelige, Schwinger:1951} one-loop effective Lagrangian in \citet{1997PhRvD..55.2449H} and expressed in terms of a set of analytic functions."
+ where and The tensors we need are derived in 7.. except that we have kept terms up to linear order in the expansion about ἐν=0 instead of dealing with the strictly wrenchless case.," where and The tensors we need are derived in \citet{1997JPhA...30.6485H}, except that we have kept terms up to linear order in the expansion about $K=0$ instead of dealing with the strictly wrenchless case."
+ Our analysis therefore requires that A« Bi}., Our analysis therefore requires that $K \ll B_k^4$ .
+, where
+Several scintillation phenomena in the interstellar medium have posed challenges to any physical model to explain them?)..,Several scintillation phenomena in the interstellar medium have posed challenges to any physical model to explain \citep{2007ASPC..365..207R}.
+ These include: (1) extreme scatering events (ESE) (?).., These include: (1) extreme scattering events (ESE) \citep{1987Natur.326..675F}.
+ Compact racio sources are occassionally observed. to. σο through a period. of demagnilicaion at low frequencies by roughly a [actor of (2) pulsar parabolic ares (2): (3) galactic center scattering., Compact radio sources are occassionally observed to go through a period of demagnification at low frequencies by roughly a factor of (2) pulsar parabolic arcs \citep{2001ApJ...549L..97S}; (3) galactic center scattering.
+ In each of these cases. a simple application of Snell's law with the assumption of spherical symmetry of the lens requires [ree electron. densities up to ~1013 5," In each of these cases, a simple application of Snell's law with the assumption of spherical symmetry of the lens requires free electron densities up to $\sim 10^4$ $^{-3}$."
+ bree electrons are at temperatures of at least  LOT. and the inferred. pressures are dillieult. to reconcile with pressure balance in the interstellar medium.," Free electrons are at temperatures of at least $\sim 10^4$ K, and the inferred pressures are difficult to reconcile with pressure balance in the interstellar medium."
+ A solution to case (3)λ has been proposed by ?.. who pointed out that scattering for sheet-like structures is dominated by the ones most aligned. with the line of sight.," A solution to case (3) has been proposed by \cite{2006ApJ...640L.159G}, who pointed out that scattering for sheet-like structures is dominated by the ones most aligned with the line of sight."
+ The alignment lowers the required three dimensional electron density., The alignment lowers the required three dimensional electron density.
+ In this paper. we compute the quantitative consequences of plasma lenses. and show —that triaxial structures are consistent with all observational data without requiring any unusual physical conditions.," In this paper, we compute the quantitative consequences of plasma lenses, and show that triaxial structures are consistent with all observational data without requiring any unusual physical conditions."
+ Large axis ratios are generic consequences of reconnection., Large axis ratios are generic consequences of reconnection.
+ In ideal resistive ΑΗ with ohmie conversion of magnetic fields. current sheets would be overdense in pressure equilibrium.," In ideal resistive MHD with ohmic conversion of magnetic fields, current sheets would be overdense in pressure equilibrium."
+ Since the resistivity. is almost certainly not ohmic. the actual density is not. known.," Since the resistivity is almost certainly not ohmic, the actual density is not known."
+ Phe phenomelogy suggests underdense current sheets. which is a probe of the physics of reconnection.," The phenomelogy suggests underdense current sheets, which is a probe of the physics of reconnection."
+ Geometric factors cause the scattering to be dominated. by ravealigned events., Geometric factors cause the scattering to be dominated by rarealigned events.
+ We follow the notation of ?.. reproducing their. lensing eeometry in reffig:lens..," We follow the notation of \cite{1992grle.book.....S}, reproducing their lensing geometry in \\ref{fig:lens}."
+ The diameter distances from the observer. to he lens plane and to the source. plane are Dy and. D. respectively. and the distance of the source plane from the ens plane is Di.," The diameter distances from the observer to the lens plane and to the source plane are ${\rm D}_{\rm d}$ and ${\rm D}_{\rm s}$ respectively, and the distance of the source plane from the lens plane is ${\rm D}_{\rm
+ ds}$."
+" Physical coordinates in the source and ens planes are i and £ respectively, defined with respect to he optic axis connecting the observer with the centre of the ens."," Physical coordinates in the source and lens planes are $\vec\eta$ and $\vec\xi$ respectively, defined with respect to the optic axis connecting the observer with the centre of the lens."
+ The deflection ofa light rav at the lens plane is denoted w à., The deflection of a light ray at the lens plane is denoted by $\vec{\hat\alpha}$.
+ The angular position of a source at iis 3. and @ is he apparent angular position from which the dellected. ray ravels.," The angular position of a source at $\vec\eta$ is $\vec\beta$, and $\vec\theta$ is the apparent angular position from which the deflected ray travels."
+" “Phev are related through the lens equation: = D.,——=—-—— theta)) =", They are related through the lens equation: = ) =
+varlalions (hal can be interpreted as the horizontal velocities (wanWerkwijkThompsonetal. 2003).,"variations that can be interpreted as the horizontal velocities \citep{vank00,tho03}."
+. In neutron stars. we expect that the amplitude of the displacements follows surface thermal variations caused by non-radial oscillations of the neutron star.," In neutron stars, we expect that the amplitude of the displacements follows surface thermal variations caused by non-radial oscillations of the neutron star."
+ As with the white dwarl stars. an increase in (he velocities can result in increased thermal emission from the star.," As with the white dwarf stars, an increase in the velocities can result in increased thermal emission from the star."
+ In our non-radial oscillation model. each frequency of oscillation has its own periodic velocily (or displacement) amplitude.," In our non-radial oscillation model, each frequency of oscillation has its own periodic velocity (or displacement) amplitude."
+ In the pulsars we have studied to date. (he amplitude of the displacements dictates (he average pulse shape (Rosen&Denmorest 2010).. but as the amplitude of the velocities grows. the average pulse shape can change (Clemens&Rosen2008).," In the pulsars we have studied to date, the amplitude of the displacements dictates the average pulse shape \citep{ros08,ros10}, but as the amplitude of the velocities grows, the average pulse shape can change \citep{cle08}."
+. In (these pulsars. through fitting a non-radial oscillation model to their cdiifting subpulses. we have also [ound (0 be large compared to that seen in white dwarf stars. consistent with their smaller size.," In these pulsars, through fitting a non-radial oscillation model to their drifting subpulses, we have also found to be large compared to that seen in white dwarf stars, consistent with their smaller size."
+ For pulsars DOSQ04-74 and BOOL3+10. values range from 126<( <133 (Rosen&Demorest2010) and 385<< «875 Clemens 2008).. respectivelv.," For pulsars B0809+74 and B0943+10, values range from $126 \leq$ $\leq 133$ \citep{ros10} and $385 \leq$ $\leq 875$ \citep{ros08}, respectively."
+ In while dwarl stars. modes with high [ade from view because of geometric cancellation of the stellar surface. leaving only modes of low 4) (Yeatesetal.2005:Thompson2004.2008).," In white dwarf stars, modes with high fade from view because of geometric cancellation of the stellar surface, leaving only modes of low $\leq 4$ ) \citep{yea05,tho04,tho08}."
+. By fitting a non-radial oscillation model to the drifting subpulses in P5ls D09434-10 and BOSO9+74. we measured the pulsation period to be on the order of 30—50 ms (Rosen& 2010).," By fitting a non-radial oscillation model to the drifting subpulses in PSRs B0943+10 and B0809+74, we measured the pulsation period to be on the order of $-$ 50 ms \citep{ros08,ros10}."
+. These values for the pulsation period are consistent with core (Reisenegeer&Goldreich1992)., These values for the pulsation period are consistent with core \citep{rei92}.
+. ILowever. core require large excitation energles. and (therefore surface are a possibility (Strohmaver1993). as they have lower energies ancl larger surface amplitudes than the core modes (McDermott 1983).," However, core require large excitation energies, and therefore surface are a possibility \citep{str93} as they have lower energies and larger surface amplitudes than the core modes \citep{mcd88}."
+. The period predictions [or core in neutron stus range from a minimum value of 10 ms (Iteisenegger&Goldreich1992) to 2-88 seconds al.1988). depending on the model for the structure and composition of (he stellar interior., The period predictions for core in neutron stars range from a minimum value of 10 ms \citep{rei92} to $-$ 88 seconds \citep{mcd88} depending on the model for the structure and composition of the stellar interior.
+ surface have periods in the 40—400 ms range (McDermottetal.1988)., Surface have periods in the $-$ 400 ms range \citep{mcd88}.
+. While the pulsation periods calculated for most models are calculated assuming low(.. the period of the mode decreases for higher spherical degree (McDermottοἱal. 1938).," While the pulsation periods calculated for most models are calculated assuming low, the period of the mode decreases for higher spherical degree \citep{mcd88}. ."
+. It is unlikely that the pulsation modes we see in PSRs D09432-10 ancl DO3004-14 are as these modes have periods on the order of tenths of milliseconds (McDermottetal.1988). and have overtones with shorter periods. whereas the overtones of have longer periods.," It is unlikely that the pulsation modes we see in PSRs B0943+10 and B0809+74 are as these modes have periods on the order of tenths of milliseconds \citep{mcd88} and have overtones with shorter periods, whereas the overtones of have longer periods."
+ The two oscillation driving mechanisms for white dwarl stars are the &— mechanism and convective driving., The two oscillation driving mechanisms for white dwarf stars are the $\kappa-\gamma$ mechanism and convective driving.
+ For pulsating DOVs stars like PG1159—035. characterized bv an abmosphere that lacks hydrogen but shows helium. carbon. and oxvgen absorption lines. the K—5 mechanism drives (he pulsations as (he opacity variessteeply with pressure (Starrlield 2003)..," For pulsating DOVs stars like $-$ 035, characterized by an atmosphere that lacks hydrogen but shows helium, carbon, and oxygen absorption lines, the $\kappa-\gamma$ mechanism drives the pulsations as the opacity variessteeply with pressure \citep{sta83,cor06,cox03}. ."
+ In DAV and DDV stars. which have atmospheres," In DAV and DBV stars, which have atmospheres"
+presence of background matter from which another nucleus can be captured (for digestion or fusion with the remainder of the old one) this time depends to some extent on the nuclear properties of the surrounding matter.,presence of background matter from which another nucleus can be captured (for digestion or fusion with the remainder of the old one) this time depends to some extent on the nuclear properties of the surrounding matter.
+ While the first. postulates are based. on the relations known from atomic and nuclear physics. the latter (ο) postulate follows already. from our experiments.," While the first postulates are based on the relations known from atomic and nuclear physics, the latter (g) postulate follows already from our experiments."
+ The discovery ofa nuclear-active radiation propagating with a velocity on the kmescale. Le. ofa population captured into particularly gcocentric orbits. was somewhat unexpected for us.," The discovery of a nuclear-active radiation propagating with a velocity on the km-scale, i.e., of a population captured into particularly geocentric orbits, was somewhat unexpected for us."
+ Therefore the next natural step was to increase the observation time of the signals from the bottom scintillator surface from LOO<Af|100pis (Drobyshevski 2000€) to 250xA:|ps with respect to the signal from the top scintillator.," Therefore the next natural step was to increase the observation time of the signals from the bottom scintillator surface from $-100 \le \Delta t \le +100\,\mks$ (Drobyshevski 2000c) to $-250 \le \Delta t \le +250\,\mks$ with respect to the signal from the top scintillator."
+ We hoped that bevond the 3100ps interval we would enter the zone of purely random signals with a No(Af)=const distribution.," We hoped that beyond the $\pm
+100\,\mks$ interval we would enter the zone of purely random signals with a $N_2(\Delta t) = {\rm const}$ distribution."
+ However this conjecture did not gain a support., However this conjecture did not gain a support.
+ We start our consideration with Af<0. Le.. with signals belonging to an upward [lux of non-monoenergetic (see Fie.," We start our consideration with $\Delta t < 0$, i.e., with signals belonging to an upward flux of non-monoenergetic (see Fig."
+ 1) particles., 1) particles.
+ Here. except a taking shape maximum at 405κα20pts. which corresponds to the 7 cm interscintillator cistance at the particle normal velocity of —23 km/s one clearly sees another three maxima. with he first one. fairky diffuse. extending from -220 to -160 us. he second. from. -120 to -100. pts. and the third. from. -80 o -60 ps.," Here, except a taking shape maximum at $-40 \le \Delta \le -20\,\mks$, which corresponds to the 7 cm interscintillator distance at the particle normal velocity of $\sim\!2-3$ km/s, one clearly sees another three maxima, with the first one, fairly diffuse, extending from -220 to -160 $\mks$, the second from -120 to -100 $\mks$, and the third, from -80 to -60 $\mks$."
+ Bearing in mind our assumptions (see Sec.2). the irst maximum can be due to the bottom scintillator excited ww the Auger electrons. which are emitted in the capture by daemons of air nuclei after getting freed of the Fe and Sn 1eavy nucle: captured in the traversal of the lower part ofthe inned-sheet casing at various angles.," Bearing in mind our assumptions (see Sec.2), the first maximum can be due to the bottom scintillator excited by the Auger electrons, which are emitted in the capture by daemons of air nuclei after getting freed of the Fe and Sn heavy nuclei captured in the traversal of the lower part of the tinned-sheet casing at various angles."
+ Phe Auger electrons can be emitted. also in the capture of Sn from the thin tin coating of the iron sheet., The Auger electrons can be emitted also in the capture of Sn from the thin tin coating of the iron sheet.
+ After the participation in all these oocesses. part of the daemons reach the top Iuminophor. and it is here that they generate the trigger pulse correlating with the signals of the first dilfuse maximum.," After the participation in all these processes, part of the daemons reach the top luminophor, and it is here that they generate the trigger pulse correlating with the signals of the first diffuse maximum."
+" The sharper second and third. maxima are possibly produced. as the daemon captures Ag (Z,= 47) and Zn (Z,= 30) nuclei. respectively. in traversing the bottom Iuminophor."," The sharper second and third maxima are possibly produced as the daemon captures Ag $Z_{\rm n} = 47$ ) and Zn $Z_{\rm n} = 30$ ) nuclei, respectively, in traversing the bottom luminophor."
+ Note that the captured S nucleus has a high probability to be shortly. lost when it encounters a Zn nucleus. anc the probability of capturing Ag is not smal despite the low concentration of Ag. because for a velocity of ~2 km/s the daemon free path until meeting an Ag atom is comparable with the ZnSXe) grain size.," Note that the captured S nucleus has a high probability to be shortly lost when it encounters a Zn nucleus, and the probability of capturing Ag is not small despite the low concentration of Ag, because for a velocity of $\sim$ 2 km/s the daemon free path until meeting an Ag atom is comparable with the ZnS(Ag) grain size."
+ On traversing a any angle the bottom polystyrene plate into the space before the top Iuminophor. daemons finish digesting the Ae aix Zn nuclei. so that when they capture air atoms. they emi here Auger electrons. which trigger the top Iuminophor.," On traversing at any angle the bottom polystyrene plate into the space before the top luminophor, daemons finish digesting the Ag and Zn nuclei, so that when they capture air atoms, they emit here Auger electrons, which trigger the top luminophor."
+ Lf our scenarium is true. the time shift between these maxima relative to Af=0 provides an idea about the time per decay of a dacmon-containing proton.," If our scenarium is true, the time shift between these maxima relative to $\Delta t = 0$ provides an idea about the time per decay of a daemon-containing proton."
+ It turns out to be IAN8ToxfyΖανf24. which shows that our estimates of τι based on Solar energeties. (Drobyvsheyski 2000a.b) were clearly unclerestimates.," It turns out to be $\Delta\tau_{\rm ex} \approx \tau_{\rm ex} / n_{\rm p} =
+\tau_{\rm ex}/Z_{\rm n} \approx 2\,\mks$, which shows that our estimates of $\tau_{\rm ex}$ based on Solar energetics (Drobyshevski 2000a,b) were clearly underestimates."
+ The position of the first maximum (GMz1ο.200 pts) vields ~1.5 km/s for the velocity with which a daemon passes normally the distance of 29 em from the lower box cover to the upper Iuminophor laver.," The position of the first maximum $\Delta t \approx 180-200\,
+\mks$ ) yields $\sim$ 1.5 km/s for the velocity with which a daemon passes normally the distance of 29 cm from the lower box cover to the upper luminophor layer."
+ Taking into account the random angular clistribution of the trajectories. the average velocity increases to 73 kms. For Al0. the daemons move downward.," Taking into account the random angular distribution of the trajectories, the average velocity increases to $\sim$ 3 km/s. For $\Delta t > 0$, the daemons move downward."
+ The first distinct maximum at ~20<Af«dps corresponds apparently to the time needed for a daemon to cross the 7-cm gap separating the scintillator lavers., The first distinct maximum at $\sim\!20 < \Delta t < 40 \mks$ corresponds apparently to the time needed for a daemon to cross the 7-cm gap separating the scintillator layers.
+ For an arbitrary trajectory inclination. this time. às we saw before (Drobvshevski 2000€). corresponds to a velocity ~3.5 kms. Le. to objects in geocentric orbits with a perigee inside the Earth.," For an arbitrary trajectory inclination, this time, as we saw before (Drobyshevski 2000c), corresponds to a velocity $\sim\!3-5$ km/s, i.e., to objects in geocentric orbits with a perigee inside the Earth."
+ Phe interpretation of this maximum. as well as the 40<A/20tts maximum. should be approached with a certain caution. because the fairly. slow digestion of nuclei by a daemon could give rise to a certain bias toward inclined. trajectories.," The interpretation of this maximum, as well as the $-40<\Delta
+t<-20\,\mks$ maximum, should be approached with a certain caution, because the fairly slow digestion of nuclei by a daemon could give rise to a certain bias toward inclined trajectories."
+ daemon poisoned by a poorly digested heavy nucleus is less likely to capture a Zn or Ag nucleus from the next scintillator laver and to initiate a scintillation init if it propagates by the fastest (direct) path., A daemon poisoned by a poorly digested heavy nucleus is less likely to capture a Zn or Ag nucleus from the next scintillator layer and to initiate a scintillation in it if it propagates by the fastest (direct) path.
+ Another selection arises because of the higher probability of capture of Ag. despite its low content. by a slower-moving daemon (the capture cross section Is proportional to V. 7).," Another selection arises because of the higher probability of capture of Ag, despite its low content, by a slower-moving daemon (the capture cross section is proportional to $V^{-2}$ )."
+ Because of the ZnSCXg) coating being not continuous. daemons with nuclei captured both in the upper scintillator have a fairly high probability of turning up also in the space under the bottom scintillator.," Because of the ZnS(Ag) coating being not continuous, daemons with nuclei captured both in the upper scintillator have a fairly high probability of turning up also in the space under the bottom scintillator."
+ Pherefore in order to reliably separate and subsequently. identifv the maxima according to their origin (the end of digestion of a Zn or Ag nucleus captured in the top or bottom scintillator ete.), Therefore in order to reliably separate and subsequently identify the maxima according to their origin (the end of digestion of a Zn or Ag nucleus captured in the top or bottom scintillator etc.)
+ for, for
+the mixed configuration,the mixed configuration
+for 6 DOF).,for 6 DOF).
+ The fitted value of us about 40 times higher than the average value of the absorption along the line of sight. which indicates that IGR JO9523-6231 1s intrinsically absorbed.," The fitted value of is about 40 times higher than the average value of the absorption along the line of sight, which indicates that IGR $-$ 6231 is intrinsically absorbed."
+ The detection of the source in the U-filter and the high X-ray absorption IGR J10147—6354 was first reported by Birdetal.(2007)., The detection of the source in the U-filter and the high X-ray absorption IGR $-$ 6354 was first reported by \citet{bird07}.
+ The source was detected in the 40-100 keV energy range at a significance of 4.9c during 1340 ks of observation., The source was detected in the 40–100 keV energy range at a significance of $\sigma$ during 1340 ks of observation.
+ The best XRT position is well within the ISGRI error box of There is a single source in the 2MASS point source catalogue within ffrom the XRT position., The best XRT position is well within the ISGRI error box of There is a single source in the 2MASS point source catalogue within from the XRT position.
+ 2MASS JI0141554-6351500 is at aaway from the centre of the error box., 2MASS J10141554-6351500 is at away from the centre of the error box.
+ It is also well detected in the UVOT UVW? filter., It is also well detected in the UVOT UVW2 filter.
+ The XRT spectrum has 93 ets for a total exposure of 4538 s. An absorbed power law fits the data well. with a C-statistic value=32.6 for 29 bins.," The XRT spectrum has 93 cts for a total exposure of 4538 s. An absorbed power law fits the data well, with a C-statistic value=32.6 for 29 bins."
+ The value of iis a factor about 7.3 times higher than the average value of the absorption along the line of sight which indicates that the absorption is mostly intrinsic to the source., The value of is a factor about 7.3 times higher than the average value of the absorption along the line of sight which indicates that the absorption is mostly intrinsic to the source.
+ The detection of a bright UV counterpart. as for the former source. suggests either a very close Galactic or a far and bright extragalactic object.," The detection of a bright UV counterpart, as for the former source, suggests either a very close Galactic or a far and bright extragalactic object."
+ Again. the low Galactic latitude would rather tend to point towards a Galactic source rather than an extragalactic one.," Again, the low Galactic latitude would rather tend to point towards a Galactic source rather than an extragalactic one."
+Iu atypical imaging observation. a considerable fraction. ifnot most. of tlie detected sources are interlopers.,"In a typical imaging observation, a considerable fraction, if not most, of the detected sources are interlopers."
+ Iu a high Galactic latitude field. interlopers are typically background Αλ». plus a small number of foregrouud stars.," In a high Galactic latitude field, interlopers are typically background AGNs, plus a small number of foreground stars."
+ The number-Ilux relation of such interlopers lias beeu characterized in various X-ray surveys. including the ACIS-I deep surveys (Moretti et al.," The number-flux relation of such interlopers has been characterized in various X-ray surveys, including the ACIS-I deep surveys (Moretti et al."
+ 2003 aud references therein)., 2003 and references therein).
+ In almost all recent work. this characterization is doue separately in the 0.5-2 keV aud 2-10 keV bands.," In almost all recent work, this characterization is done separately in the 0.5-2 keV and 2-10 keV bands."
+ The source populatious detected in these two bauds only. partially overlap: Sources with steep X-ray spectra preferentially appear in the 0.5-2 keV baud. whereas highly-absorbec oues appear in the 2-10 keV baud.," The source populations detected in these two bands only partially overlap: Sources with steep X-ray spectra preferentially appear in the 0.5-2 keV band, whereas highly-absorbed ones appear in the 2-10 keV band."
+ In general. is much more seusitive to sources iu the former baud than in the latter baud.," In general, is much more sensitive to sources in the former band than in the latter band."
+ Therefore. we focus ou the accumulated nuimnber-flux relation in the 0.5-2 keV band. whieh can be modeled as: where Nj=6150sourcesdeg7. αι= 1.82. ao= 0.60.and Sy=1.18x10.therestem7 ," Therefore, we focus on the accumulated number-flux relation in the 0.5-2 keV band, which can be modeled as: where $N_0 = 6150 {\rm~sources~deg^{-2}}$ $\alpha_1 = 1.82$ , $\alpha_2 = 0.60$ ,and $S_0 = 1.48 \times 10^{-14} {\rm~erg~s^{-1}~cm^{-2}}$ "
+Iu the right-hand side o this equation the first term in brackets describes the 1‘clative role of D Inning iu the total energy budget. second teiu represeuts intrinsic stellar luminosity. while tie third term is due to the inflow of the eravitational potential enerev GALAL/R the ratio of this cucrey infow to Ly differs from A only by a coustant factor.,"In the right-hand side of this equation the first term in brackets describes the relative role of $D$ burning in the total energy budget, second term represents intrinsic stellar luminosity, while the third term is due to the inflow of the gravitational potential energy $GM\dot M/R$ – the ratio of this energy inflow to $L_0$ differs from $\Lambda$ only by a constant factor."
+ Suppression of L is hugest when A is highest which Is casy fo see by inspecing Figures 1—1.., Suppression of $L$ is largest when $\Lambda$ is highest which is easy to see by inspecting Figures \ref{fig:f1}- \ref{fig:f4}.
+ But this automatically meaus that the maxima deviation of the suppression factor X from unitv occus preciselv when the inflow of the gravitational potential energv far exceeds the stella lununosity., But this automatically means that the maximum deviation of the suppression factor $\chi$ from unity occurs precisely when the inflow of the gravitational potential energy far exceeds the stellar luminosity.
+ Apparently. uuder these circumstances £L is a subdominant contribution to the stellar energv budect and thus even a significant reduction of L compared to Ly is going to be uceligible conrpared to the eravitational energy influx.," Apparently, under these circumstances $L$ is a subdominant contribution to the stellar energy budget and thus even a significant reduction of $L$ compared to $L_0$ is going to be negligible compared to the gravitational energy influx."
+ Moreover. Figure 1 deuoustrates that A reaches its asin when AB> is at its nüunauunu while fp is still very close to unitv.," Moreover, Figure \ref{fig:f1} demonstrates that $\Lambda$ reaches its maximum when $R$ is at its minimum while $f_D$ is still very close to unity."
+ At fus point vigorous D burnius columences Inside the sar giving rise to wvorv hieh LpíLg., At this point vigorous D burning commences inside the star giving rise to very high $L_D/L_0$.
+ As a result. at tie evolutionary stage when \ is nuninial £ is ποταιant in comparison to not ouly GALAL/R but also Lp.," As a result, at the evolutionary stage when $\chi$ is minimal $L$ is subdominant in comparison to not only $GM\dot M/R$ but also $L_D$."
+ Tus additionally downplays the role of the hunuinositv suppression by iradiatiou iu the carly protostellay evolution., This additionally downplays the role of the luminosity suppression by irradiation in the early protostellar evolution.
+ This line of reasoning aso explaius why at AM~1 M. we have found AR/R to )o larger for lower AL (see 83)]., This line of reasoning also explains why at $M\sim 1$ $_\odot$ we have found $\Delta R/R$ to be larger for lower $\dot M$ (see \ref{sect:res}) ).
+ First. analler M. ποστς lower X so that the ratio of L to the gravitational energy 1vow rate CAZAZR in the low AT case is larecr than iu ιο hieh AL case.," First, smaller $\dot M$ means lower $\Lambda$ so that the ratio of $L$ to the gravitational energy inflow rate $GM\dot M/R$ in the low $\dot M$ case is larger than in the high $\dot M$ case."
+ Also. at AL~ AL. one generally finds f;»l (see Figure lee) so that Lp isiaiuly due to the birue of the freshly accreted D (rather tman the D that remained iu the protostar from previous accretion).," Also, at $M\sim$ $_\odot$ one generally finds $f_D\ll 1$ (see Figure \ref{fig:f1}e e) so that $L_D$ is mainly due to the burning of the freshly accreted D (rather than the D that remained in the protostar from previous accretion)."
+ Since in the low AL case less fresh D is supplied to the proostar Lp ust also be lower than iu he high AL case., Since in the low $\dot M$ case less fresh $D$ is supplied to the protostar $L_D$ must also be lower than in the high $\dot M$ case.
+ As a result. iu the lower AL case L plavs a iore significant role compared to Ly (iu which case ARR should © niore sensitive to changes iu L cause by radiation) thau iu the hieh AT case.," As a result, in the lower $\dot M$ case $L$ plays a more significant role compared to $L_D$ (in which case $\Delta R/R$ should be more sensitive to changes in $L$ caused by irradiation) than in the high $\dot M$ case."
+ This conclusion munediately raises the following question: since AR/R ineoases as AL decreases would ouc find AR/R~1 at knv enough Jf?, This conclusion immediately raises the following question: since $\Delta R/R$ increases as $\dot M$ decreases would one find $\Delta R/R\sim 1$ at low enough $\dot M$?
+ The answer is no. and it has to do with the fact that 4 appreciadv differs from unity (obviously. a necessary couditiou for ectting AR/R~ 1) oul~ at rather high A.," The answer is no, and it has to do with the fact that $\chi$ appreciably differs from unity (obviously, a necessary condition for getting $\Delta R/R\sim 1$ ) only at rather high $\Lambda$ ."
+ This is a eeneric feature of disk irradiation which is illustraed in Fieure 5 where we display Aso the value of X at which \(Asu)=15 as a function of assmued opacity law represented by the parameter ©. see equation (11).," This is a generic feature of disk irradiation which is illustrated in Figure \ref{fig:f5} where we display $\Lambda_{50}$ – the value of $\Lambda$ at which $\chi(\Lambda_{50})=0.5$ – as a function of assumed opacity law represented by the parameter $\xi$, see equation \ref{eq:xi}) )."
+ One can see that Asy2 aD↓∩−↕∪↥⋅���↧↕≦↓∙⋯↸∖⋜↧⋯∐∶↴⋁ ⋅ that significant hlunimositv suppression requires rather lueh AJ.," One can see that $\Lambda_{50}\gtrsim 10^2$ for all $\xi>4$, meaning that significant luminosity suppression requires rather high $\dot M$."
+ This meffücienev of imradialon 1n suppressing Lis caused by the specific eeonetrv of disk inadiatiou in which the radiation flux is a very seusitive fiction (x 0°) of the latitude at t1je stellar surface 0. sec Rafikov (2007).," This inefficiency of irradiation in suppressing $L$ is caused by the specific geometry of disk irradiation in which the irradiation flux is a very sensitive function $\propto \theta^5$ ) of the latitude at the stellar surface $\theta$, see Rafikov (2007)."
+ Because of that stellar polar caps can stay cool even at rather high AY allowing unsuppressed flux to be cluitted over a significant portion of the stellar surface., Because of that stellar polar caps can stay cool even at rather high $\dot M$ allowing unsuppressed flux to be emitted over a significant portion of the stellar surface.
+ Iu the case of &£=6.5 as appropriate for cool. OW-1uass protostars one finds that Asy=2.2« which according to equation (9)) muuediatelv iuplics hat GALM/IRxL75Ly when 4=0.5.," In the case of $\xi=6.5$ as appropriate for cool, low-mass protostars one finds that $\Lambda_{50}=2.2\times 10^3$ which according to equation \ref{eq:Lambda}) ) immediately implies that $GM\dot M/R\approx 175 L_0$ when $\chi=0.5$."
+ Clearly. in his case stellar huuinositv should have μια] effect ou he xotostellu evolution.," Clearly, in this case stellar luminosity should have small effect on the protostellar evolution."
+ If AJ ds so that GALAL/R~Ly (aud A~ 1) stellar hunuinositv would ος plaving a siguificaut role in the stellar energv budeet. rowever x would be very close to unity (see Rafikov 20Nn) and the £ suppression by irradiation wold be negligible.," If $\dot M$ is so that $GM\dot M/R\sim L_0$ (and $\Lambda\sim 1$ ) stellar luminosity would be playing a significant role in the stellar energy budget, however $\chi$ would be very close to unity (see Rafikov 2007) and the $L$ suppression by irradiation would be negligible."
+ Thus. 1udder no circiustauces should oue expect AR/R argertlan several per celt. camethat quite ecucrally he iradiation. by accretion disk is uulikelv to play a significant role iu the evolution of the protostellar xoperties.," Thus, under no circumstances should one expect $\Delta R/R$ larger than several per cent, meaningthat quite generally the irradiation by accretion disk is unlikely to play a significant role in the evolution of the protostellar properties."
+role here because there is no rotational component to the How.,role here because there is no rotational component to the flow.
+ Figure 6 summarizes the hydrodynamie state variables at time /=1l., Figure \ref{fig:Sod-200} summarizes the hydrodynamic state variables at time $t=1$.
+ To first approximation one ects identical results with the new formalism as compared to the standard approach., To first approximation one gets identical results with the new formalism as compared to the standard approach.
+ Linear momentum is not conserved inrpSPtl and we find a linear excess velocity of (105.3.10.7) so per particle an error on the velocity of 0.003 in the a and a completely negligible component along the gy-direction.," Linear momentum is not conserved in and we find a linear excess velocity of $(-105,-3\times 10^{-5})$ so per particle an error on the velocity of $0.003$ in the $x$ -direction and a completely negligible component along the $y$ -direction."
+ This is at a time when the r.nis., This is at a time when the r.m.s.
+ velocity is O46 so just slightly. above one half of a per cent. error in the dominant w-velocity., velocity is $\sim 0.46$ so just slightly above one half of a per cent error in the dominant $x$ -velocity.
+ SPILL has poor behaviour at the contact discontinuities., SPH has poor behaviour at the contact discontinuities.
+ For both the one originating [roni he initially smoothec and the the discontinuous interface at the right boundary., For both the one originating from the initially smoothed and the the discontinuous interface at the right boundary.
+ Both contacts at ik~3.8 and μον9.3 are better captured byΡΟΗ., Both contacts at $x\sim3.8$ and $x\sim 9.3$ are better captured by.
+ The Sod shock tube has few catures and it is reassuring that using as many as 2007 particles can give an excellent answer., The Sod shock tube has few features and it is reassuring that using as many as $200^2$ particles can give an excellent answer.
+ There are only slight dilferences in howrpSPH handles one dimensional shock tubes., There are only slight differences in how handles one dimensional shock tubes.
+ We will discuss one very »opular application taken from a cosmological context after esting a very strong shock next., We will discuss one very popular application taken from a cosmological context after testing a very strong shock next.
+ llere we give another test of a much stronger shock than the one by sod., Here we give another test of a much stronger shock than the one by sod.
+ This one has a Mach number close to one hundred., This one has a Mach number close to one hundred.
+ We also use the chance to compare this to the cillerence formulation studied by 2.., We also use the chance to compare this to the difference formulation studied by \cite{1996PASA...13...97M}.
+ Phe density. and pressure are (1.6.61033 on the left and (1/5.1) on the right.," The density and pressure are $(1,6.6\times
+10^4)$ on the left and $(1/5,1)$ on the right."
+ ‘This is very similar to the one studied by ? and is well known to work well with standard SPIL, This is very similar to the one studied by \cite{2006MNRAS.367..113P} and is well known to work well with standard SPH.
+L Leere we use 35 neighbours. a—4 and 5000 particles.," Here we use $35$ neighbours, $\alpha=4$ and 5000 particles."
+ This is à good example where one can makerpSPH and the Morris formulation give unphysical results., This is a good example where one can make and the Morris formulation give unphysical results.
+ These methods require the pressure gradient to be resolved., These methods require the pressure gradient to be resolved.
+ So if vou start with completely discontinuous left right states one will eet unphysical waves giving unexpected results., So if you start with completely discontinuous left right states one will get unphysical waves giving unexpected results.
+ However. this is not a shortcoming of the method but simply are errors that come from not resolving the initial conditions.," However, this is not a shortcoming of the method but simply are errors that come from not resolving the initial conditions."
+ We again use the ramp function from above with a width of4 in this very long domain ranging from 0 to 500 in.r and Oto 10 in y., We again use the ramp function from above with a width of 4 in this very long domain ranging from 0 to 500 in $x$ and $0$ to $10$ in $y$.
+ We cannot confirm Morris! claim that his formulation gives large post shock oscillations in this method and suspect that he may have set up discontinuous initial conditions., We cannot confirm Morris' claim that his formulation gives large post shock oscillations in this method and suspect that he may have set up discontinuous initial conditions.
+ We can see that our new formulation performs somewhat better than the Morris formulation as it does not overshoot the analvtical density jump of 4 raising the density from 0.2 to 0.8 in Figure 7.., We can see that our new formulation performs somewhat better than the Morris formulation as it does not overshoot the analytical density jump of 4 raising the density from $0.2$ to $0.8$ in Figure \ref{fig:StrongShock}.
+ Otherwise both approaches work fine and have no problem in modeling strong shocks and evolving it for large distances., Otherwise both approaches work fine and have no problem in modeling strong shocks and evolving it for large distances.
+ Another particularly strong shock is formed in the Sedov-‘Tavlor blast wave (2). presenting a cdillieult test. problem for incompressible hvdrodynamies codes., Another particularly strong shock is formed in the Sedov-Taylor blast wave \citep{1959flme.book.....L} presenting a difficult test problem for incompressible hydrodynamics codes.
+ One the one hand it is a self similar solution which makes it insensitive to how exactly one sets it up as long as one evolves the svsten for a very long time., One the one hand it is a self similar solution which makes it insensitive to how exactly one sets it up as long as one evolves the system for a very long time.
+ On the other hand it is the solution for a point explosion., On the other hand it is the solution for a point explosion.
+ For a given finite resolution. however. there is no unique wav of specifving the initial conditions.," For a given finite resolution, however, there is no unique way of specifying the initial conditions."
+ llere is where exact momentum anc energy. conservation is very helpful as one can set. up the initial conditions at will and even if one were to make verv large errors in the time evolution the method will still arrive at. the self similar solution., Here is where exact momentum and energy conservation is very helpful as one can set up the initial conditions at will and even if one were to make very large errors in the time evolution the method will still arrive at the self similar solution.
+ In conservative grid codes this still can [σας to aspherical solutions if one did. not resolve the spherical central hot region., In conservative grid codes this still can lead to aspherical solutions if one did not resolve the spherical central hot region.
+ Since SPLL however uses spherical kernels one can get away sometimes even by just heating one single particle (2)..., Since SPH however uses spherical kernels one can get away sometimes even by just heating one single particle \citep{2002MNRAS.333..649S}.
+ This is very useful in applications such as ealaxy formation simulations where one is always far from resolving the relevant length scales of an explosion., This is very useful in applications such as galaxy formation simulations where one is always far from resolving the relevant length scales of an explosion.
+ On the other hand any physies that were to occur at a scale. of the shell thickness would. be impossible to resolve in such a single particle energy ejection., On the other hand any physics that were to occur at a scale of the shell thickness would be impossible to resolve in such a single particle energy ejection.
+ ForSPL and the Morris formulation we need to resolve the pressure gradients in the initial conditions as we saw above in the strong shock setup., For and the Morris formulation we need to resolve the pressure gradients in the initial conditions as we saw above in the strong shock setup.
+ We setup a square lattice of particles with 300 particles on à side in the unit domain., We setup a square lattice of particles with 300 particles on a side in the unit domain.
+ For resolved initial conditions we set a spherical region in the center of radius r=0.1 with the same ramp function as above using a width of 0.1. to have a sound speed of one for an aciahatic index of 55/3.," For resolved initial conditions we set a spherical region in the center of radius $r=0.1$ with the same ramp function as above using a width of $0.1$, to have a sound speed of one for an adiabatic index of $\gamma=5/3$."
+ For both simulations we used a Courant number of 0.2 (0.1 in Gadget). SO neighbors. artificial viscosity à=2.5 and had the Balsara switch olf.," For both simulations we used a Courant number of 0.2 $0.1$ in Gadget), 80 neighbors, artificial viscosity $\alpha=2.5$ and had the Balsara switch off."
+ Figure S shows that there potentially is also an advantage torpSPL simulations when modeling shocks., Figure \ref{fig:Sedov} shows that there potentially is also an advantage to simulations when modeling shocks.
+ We have Buled to get standard SPLE to give a stable correct density jump of (5Ες1)=4 for our setup., We have failed to get standard SPH to give a stable correct density jump of $(\gamma+1)/(\gamma-1)=4$ for our setup.
+ Also the three dimensional versions shown by ο always sccm to be too low bv as much as a factor of two., Also the three dimensional versions shown by \cite{2002MNRAS.333..649S} always seem to be too low by as much as a factor of two.
+ However. stanclare SPIL is much less sensitive to how one sets up the initia conditions and performs much better at low resolutions.," However, standard SPH is much less sensitive to how one sets up the initial conditions and performs much better at low resolutions."
+ These strong shock problems can work but clearly are not the biggest strengths ofSPL., These strong shock problems can work but clearly are not the biggest strengths of.
+ Llowever. at this poin we simply reused the artificial viscosity. prescription which was designed for standard SPL.," However, at this point we simply reused the artificial viscosity prescription which was designed for standard SPH."
+ We believe it is likely tha one can find an alternative formulation for the artificia viscosity that fits better into the discretisation which maw improve its behaviour for highly supersonic conditions., We believe it is likely that one can find an alternative formulation for the artificial viscosity that fits better into the discretisation which may improve its behaviour for highly supersonic conditions.
+ Until à better artificial viscosity prescription is designed one may opt to switch between standard SPIEL and based on the local divergence., Until a better artificial viscosity prescription is designed one may opt to switch between standard SPH and based on the local divergence.
+ We have successfully applied this strategy by using a switch that evaluates the standard SPL sum if. h;:5;73602; and the sum for less strongly convergent Bow.," We have successfully applied this strategy by using a switch that evaluates the standard SPH sum if $-h_i \div \vec{v}_i > 3
+c_{s,i}$ and the sum for less strongly convergent flow."
+ Here 5h; and ον denote the smoothing length and the current sound. speed. of the particle.," Here $h_i$ and $c_{s,i}$ denote the smoothing length and the current sound speed of the particle."
+ Εις formulation is robust in all our tests., This formulation is robust in all our tests.
+ In 1995 a comparison project was initiated. that. aimed to compare all numerical cosmology codes at that time for relevant. realistic initial conditions., In 1995 a comparison project was initiated that aimed to compare all numerical cosmology codes at that time for relevant realistic initial conditions.
+ Phe study. focused on three dimensional calculations of the formation of a ealaxy cluster in the standard. CDM. scenario of structure formation., The study focused on three dimensional calculations of the formation of a galaxy cluster in the standard CDM scenario of structure formation.
+ The choice was a setup which does not include any other physies than cosmological hyelrodvnanics with an ideal gas equation of state (often referred to as adiabatic simulations despite the entropy. generation in shocks)., The choice was a setup which does not include any other physics than cosmological hydrodynamics with an ideal gas equation of state (often referred to as adiabatic simulations despite the entropy generation in shocks).
+ The study. produced a detailed report in (7.E99. herafter).. ," The study produced a detailed report in \cite[][F99, herafter]{1999ApJ...525..554F}. ."
+One of the most surprising findings of the study was that while there was very good agreement. between the six dillerent, One of the most surprising findings of the study was that while there was very good agreement between the six different
+s!.,.
+". Neither the oscillation of the velocity, the continuum nor the core intensity seem to be related to the magnetic flux density (and area) fluctuations, as can be seen in Fig. 3.."," Neither the oscillation of the velocity, the continuum nor the core intensity seem to be related to the magnetic flux density (and area) fluctuations, as can be seen in Fig. \ref{var_resto}."
+" Whereas some of the oscillations in area we observe display periods in the range of the p-modes, others have distinctly longer periods, even up to 11 min, and show an abrupt change of the wave period at a given time."," Whereas some of the oscillations in area we observe display periods in the range of the p-modes, others have distinctly longer periods, even up to 11 min, and show an abrupt change of the wave period at a given time."
+" Moreover, we find that the area of some magnetic patches located"," Moreover, we find that the area of some magnetic patches located"
+The birth of stars. a process so conuuoulv occurring hroughout the whole universe. is as puzzling as it is ascinating.,"The birth of stars, a process so commonly occurring throughout the whole universe, is as puzzling as it is fascinating."
+" The current paradienài for star formation. ramed originally by Rant and Laplace in the 1s""! century. sugecsts that stars are born via gravitational collapse of the dense cores of molecular clouds. iu urn coudeused out of the diffuse interstellar medi."," The current paradigm for star formation, framed originally by Kant and Laplace in the $18^{\rm th}$ century, suggests that stars are born via gravitational collapse of the dense cores of molecular clouds, in turn condensed out of the diffuse interstellar medium."
+ Unufortunatelv. this compelling visualisation aud its physical uuderpimniug are iu apparent contradiction.," Unfortunately, this compelling visualisation and its physical underpinning are in apparent contradiction."
+ The reason for the above statement is that couservatioulit of angular moment curing the collapse results in he progressive inerease of the centrifugal force. which eventually ialts he iufaliug eas aud leads to the development of a central mass (16. a “protostar’) surrounded by a flattened disc of material (an accretion disc’).," The reason for the above statement is that conservation of angular momentum during the collapse results in the progressive increase of the centrifugal force, which eventually halts the infalling gas and leads to the development of a central mass (i.e. a `protostar') surrounded by a flattened disc of material (an `accretion disc')."
+ Observational evidence for the presence of such discs around voune stellar objects is compelling., Observational evidence for the presence of such discs around young stellar objects is compelling.
+ It conrprises nuaenmg iu the near infrared and optica wavebauds (6.8.seethereviewsbyWatsonetal.2007:AleCaughreanctal.2000) as well as interferometric studies that have resolved. the velocity profile ar structure in the immer regions of cliscs. up to ~ a few tens of AU from the ceutre (e.g.Dutreyetal.2007:Wilner&Lav 2000).," It comprises imaging in the near infrared and optical wavebands \citep[e.g.~see the reviews by][]{WSWM07, MSC00} as well as interferometric studies that have resolved the velocity profile and structure in the inner regions of discs, up to $\sim$ a few tens of AU from the centre \citep[e.g.][]{DGH07, WL00}."
+. This phenomenon is. in fact. not limite to voung sars.," This phenomenon is, in fact, not limited to young stars."
+ Many astroplivsical svstenis exlibit the characteristic disc-like structure that naturally results when iuw:ud notion in the plane of rotation is restricted by augular momentum conscrvation. while collapse continues in the perpendicular (polar) directiou.," Many astrophysical systems exhibit the characteristic disc-like structure that naturally results when inward motion in the plane of rotation is restricted by angular momentum conservation, while collapse continues in the perpendicular (polar) direction."
+§f Such structures are commonly associated. for example. with the «iscs of material feeding the cores of active ealaxies aux black holes.," Such structures are commonly associated, for example, with the discs of material feeding the cores of active galaxies and black holes."
+ The difficulty iu. progressing past this stage ds clear when we recall that these protostellar discs are dynamically stable., The difficulty in progressing past this stage is clear when we recall that these protostellar discs are dynamically stable.
+ Ao typica disc is cifferentiallv rotating. with an aneular monenutun (L) profile that," A typical disc is differentially rotating, with an angular momentum $L$ ) profile that"
+"A series of exterior MMRs with Neptune shows strong peaks at all grain sizes in the Tmaxv1077 model: 4:3 3:2 7:4 2:1 (2236, 39, 44 and 48 AU).","A series of exterior MMRs with Neptune shows strong peaks at all grain sizes in the $\tau_{\rm{max}} \sim 10^{-7}$ model: 4:3 3:2 7:4 2:1 $\approx$ 36, 39, 44 and 48 AU)."
+ Dashed lines show these and a few other MMRs in Figure 4.., Dashed lines show these and a few other MMRs in Figure \ref{fig:semimajoraxis1}.
+ Some of these peaks also survive in the total semimajor axis distribution., Some of these peaks also survive in the total semimajor axis distribution.
+" They appear in all three source populations, though they are strongest in the cold classical population, probably these objects have small eccentricities and inclinations, making them easier to trap in MMRs."," They appear in all three source populations, though they are strongest in the cold classical population, probably these objects have small eccentricities and inclinations, making them easier to trap in MMRs."
+" Though the plutino dust is released from bodies in Neptune's 3:2 MMR, only grains larger than 19.7 wm remain tightly concentrated around that MMR, at z39 AU."," Though the plutino dust is released from bodies in Neptune's 3:2 MMR, only grains larger than $19.7 \ \mu$ m remain tightly concentrated around that MMR, at $\approx 39$ AU."
+" Figure 3 shows that at this dust level (Tax~ 1077), the resonant"," Figure \ref{fig:collisions} shows that at this dust level $\tau_{\rm{max}} \sim 10^{-7}$ ), the resonant"
+Table 3 shows the resulting fractions for comparison of the distributions of detected parameters and one year of observation.,Table \ref{tab:comptable_1yr} shows the resulting fractions for comparison of the distributions of detected parameters and one year of observation.
+ All of the models can be consistently distinguished based on some combination of one or two of their parameter distributions: Naturally. these numbers improve when a longer observation window affords more sources to make the comparison.," All of the models can be consistently distinguished based on some combination of one or two of their parameter distributions: Naturally, these numbers improve when a longer observation window affords more sources to make the comparison."
+ Table 4 shows comparison results when sources are chosen from equation 8 according to à 3 year observation window., Table \ref{tab:comptable_3yr} shows comparison results when sources are chosen from equation \ref{psprob} according to a 3 year observation window.
+ With the 3 years of sources. the only model parameter distributions that cannot be reliably distinguished are the SE vs. SC distributions and the LE vs. LC distributions. as functions of the masses and 25.," With the 3 years of sources, the only model parameter distributions that cannot be reliably distinguished are the SE vs. SC distributions and the LE vs. LC distributions, as functions of the masses and $D_L$."
+" The SE and SC distributions of masses and D,. in particular. are not generally distinguishable."," The SE and SC distributions of masses and $D_L$, in particular, are not generally distinguishable."
+" The distributions of InD; and Iuij, for the LE and LC models. on the other hand. are often distinguishable."," The distributions of $\ln{D_L}$ and $\ln{m_1}$ for the LE and LC models, on the other hand, are often distinguishable."
+" Since we are interested as much in the limits of LISA's ability to distinguish between the population models as in whether or not LISA will be able to distinguish between these particular sets of models (which will undoubtedly be superseded with improved versions by the time LISA is operational). Figures 20. and 21. show the distributions of InD, and Inn. and of InD; and Ini)»."," Since we are interested as much in the limits of LISA's ability to distinguish between the population models as in whether or not LISA will be able to distinguish between these particular sets of models (which will undoubtedly be superseded with improved versions by the time LISA is operational), Figures \ref{fig:ExampleDLm1distributions} and \ref{fig:ExampleDLm2distributions} show the distributions of $\ln{D_L}$ and $\ln{m_1}$, and of $\ln{D_L}$ and $\ln{m_2}$."
+ The plots serve to illustrate distributions which have varying levels of distinguishability., The plots serve to illustrate distributions which have varying levels of distinguishability.
+ In order to assess the significance of the effects of parameter estimation error on the distinguishability of the models. we have also performed model comparison tests with no parameter estimation errors.," In order to assess the significance of the effects of parameter estimation error on the distinguishability of the models, we have also performed model comparison tests with no parameter estimation errors."
+ The results of these tests. for one year of observations. are shown in Table5..," The results of these tests, for one year of observations, are shown in Table\ref{tab:comptable_noerrs}."
+ The differences between these model distinguishability fractions and those found in Table 3 (which incorporate the parameter estimation errors) are shown in Table 6.., The differences between these model distinguishability fractions and those found in Table \ref{tab:comptable_1yr} (which incorporate the parameter estimation errors) are shown in Table \ref{tab:comptable_differences}. .
+ These, These
+The relation between the skv map we seek and the observed data stream may be cast as a linear algebra system (Wrightetal1996:Teemark1997).,"The relation between the sky map we seek and the observed data stream may be cast as a linear algebra system \cite{WrHi96a,Te97}."
+ Let ; and p Judices denote quantities iu the temporal auc spatial domains. and eroup as a data vector. dy. and a noise vector the temporal stream of collected data aud the detector noise stream. both of dimension Ay.," Let $_t$ and $_p$ indices denote quantities in the temporal and spatial domains, and group as a data vector, $d_{t}$, and a noise vector the temporal stream of collected data and the detector noise stream, both of dimension $\mathcal{N}_{tod}$ ."
+ We then have |Hg. where obey ds the sigual vector given bw the observation of the uuknown pixelised sky map. or). which lasbeen arranged as a vector of dimecusion Nui," We then have = +, where $A_{tp}x_{p}$ is the signal vector given by the observation of the unknown pixelised sky map, $x_{p}$, which hasbeen arranged as a vector of dimension $\mathcal{N}_{pix}$."
+" The Nou- ""observation"" matrix 4 therefore encompasses hne scanunimg strategv aud the beam pattern of the detector."," The $\mathcal{N}_{tod} \times 
+\mathcal{N}_{pix}$ “observation” matrix $A$ therefore encompasses the scanning strategy and the beam pattern of the detector."
+ Iu the following. we restrict to the case when the beam ρατοι ds svnuuetrical.," In the following, we restrict to the case when the beam pattern is symmetrical."
+" We ean therefore take wr, to be a nap of the skv which has already been couvolved. with he beau pattern. and <A oulv deseribes how this sky is reine scanned."," We can therefore take $x_{p}$ to be a map of the sky which has already been convolved with the beam pattern, and $A$ only describes how this sky is being scanned."
+ For the total power iieasurement nou-ifferential mieasurenient) we are iterested im here. the joservation matrix A then has a single non-zero clement scr row. Which can be set to one if d and we are expressed oe ithe same units.," For the total power measurement non-differential measurement) we are interested in here, the observation matrix $A$ then has a single non-zero element per row, which can be set to one if $d$ and $x$ are expressed in the same units."
+ The model of the measurement is then ute transparent: cach temporal datum is the simu of he pixel value to which the detector is pointing plus the ctector nolse., The model of the measurement is then quite transparent: each temporal datum is the sum of the pixel value to which the detector is pointing plus the detector noise.
+ The map-makine step then amounts to best solve for .e eiven d (aud some properties of the noise)., The map-making step then amounts to best solve for $x$ given $d$ (and some properties of the noise).
+" We shall seek a estimator of .,,. Wed;e "," We shall seek a estimator of $x_{p}$, = ."
+To motivate a particular choice of the NySN matrix VW. a Bavesiau approach is couvenicut.," To motivate a particular choice of the $\mathcal{N}_{pix} \times 
+\mathcal{N}_{tod}$ matrix $W$, a Bayesian approach is convenient."
+" Iudeed we are secking the optimal solution to this inversion problem which maximises the probability of a deduced set of theory paralucters (here the map .,}) given our data (d;) by inaxinising Pr|d).", Indeed we are seeking the optimal solution to this inversion problem which maximises the probability of a deduced set of theory parameters (here the map $x_p$ ) given our data $d_t$ ) by maximising $\mathcal{P}(x|d)$.
+ Baves’ thieoreii simply states that Pat =PidPld)., Bayes' theorem simply states that (x|d) =.
+ If weprior. thon 0 follows a uniform distribution as well as 4.," If we, then $x$ follows a uniform distribution as well as $d$."
+ Therefore. Pusd) Plait)If we further assume thatCatssian.. we cal write Pitsd) erpl 215} ορ(dAryeNIU where Ny=<?με is the noise covariance matrix.," Therefore, (x|d) (d|x).If we further assume that, we can write (x|d) exp(- /2) exp(-(d-Ax)_t^T where $N_{tt'} = <nn^T>_{tt'}$ is the noise covariance matrix."
+" Iu this particular ease. maximising PCr|y) amounts to find the least square solution which was used to analyse the ""CODIZ data (JansenandCulkis1992).. =DATN 3. Tn this paper we will actually deal onlv with this estimator."," In this particular case, maximising $\mathcal{P}(x|y)$ amounts to find the least square solution which was used to analyse the “COBE” data \cite{JaGu92}, , W = A^T. In this paper we will actually deal only with this estimator."
+" Nevertheless as a next iteration in the analysis process, we could corporate various theoretical priors by explicitiug P(r)."," Nevertheless as a next iteration in the analysis process, we could incorporate various theoretical priors by expliciting $\mathcal{P}(x)$."
+" For example. it is often assumed for the theory. Por)xccptwlpple’2) wherehore €.C,=6pd dsix the signalsienal covariiucecovariiMN matrix."," For example, it is often assumed for the theory, $\mathcal{P}(x) \propto
+exp(-x^T_pC^{-1}_{pp'}x_{p'}/2)$ where $C_{pp'}=\langle x_px^T_{p'} \rangle$ is the signal covariance matrix."
+atrix Iu that case theαμήν particular solution turus out to be the Wiener filtering solution (Zaroubietal.1995:BouchetaudCüspoer 1998): ={C cq APN TAY)1ATN 1 But this solution παν always be obtained by a further (Wiener) filtering of the CODE solution. aud we do not consider it further.," In that case the particular solution turns out to be the Wiener filtering solution \cite{ZaHo95,BoGi96,TeEf96,BoGi98}: W = + A^T A A^T. But this solution may always be obtained by a further (Wiener) filtering of the COBE solution, and we do not consider it further."
+ The prior-less solution demoustrates that as long as the (Caussian) iustrumental noise is not white. a «ΠΙΟ averaging (co-addition) of all the data poiuts corresponding oa given skv pixel is not optimal.," The prior-less solution demonstrates that as long as the (Gaussian) instrumental noise is not white, a simple averaging (co-addition) of all the data points corresponding to a given sky pixel is not optimal."
+ If the ποῖκο exhibits some temporal correlations. as induced for instance by a low-frequeucy 1/f behavior of the noise spectrin which prevails iun most CAIB experiueuts. one its to take into account the full time correlation structure of the noise.," If the noise exhibits some temporal correlations, as induced for instance by a low-frequency $1/f$ behavior of the noise spectrum which prevails in most CMB experiments, one has to take into account the full time correlation structure of the noise."
+ Additionally. this expression demonstrates hat even if the noise has a simple time structure. he scanning strateev eenerically iuduces a non-trivial correlation matrix [AENL4] ofthe noise map.," Additionally this expression demonstrates that even if the noise has a simple time structure, the scanning strategy generically induces a non-trivial correlation matrix $\left[ A^T N^{-1} A\right]^{-1}$ of the noise map."
+ Even if the xoblem is well posed formalhy. a quick looks thyorders of magnitude shows that the actual finding of i solution is non trivia task.," Even if the problem is well posed formally, a quick look at the orders of magnitude shows that the actual finding of a solution is non trivial task."
+" Iudeed a brute force method anune at inverting the full matrix [APTN14] loan operation scaliug as ΟΛ). is already. hardly tractable for present long duration balloon flights as MAXIMA. BOOMERanG. ARCIIEOPS or Topllat where A,~ and Ny 10°, "," Indeed a brute force method aiming at inverting the full matrix $\left[ A^T N^{-1} A\right]^{-1} \displaystyle$, an operation scaling as $\mathcal{O}(\mathcal{N}_{pix}^3)$, is already hardly tractable for present long duration balloon flights as MAXIMA, BOOMERanG, ARCHEOPS or TopHat where $\mathcal{N}_{tod}\sim 10^6$ and $\mathcal{N}_{pix} \sim 10^5$ ."
+It appears totally impractical for PLANCK since for a single detector (nid LOS) Noo10? and Au~ 10*!, It appears totally impractical for PLANCK since for a single detector (amid $10$ s) $\mathcal{N}_{tod}\sim 10^9$ and $\mathcal{N}_{pix} \sim 10^7$ !
+ Onepossibility may be to take advantage of specific scanniue strategies. and actually solvethe inverse of the convolution problem as detailed in," Onepossibility may be to take advantage of specific scanning strategies, and actually solvethe inverse of the convolution problem as detailed in"
+This is not trivial. as the rotation velocities and disk. Iuuimositfies are determined completely independently (and by different workers) and ouly the one using the absolute magnitude necds an assumption for the distance scale.,"This is not trivial, as the rotation velocities and disk luminosities are determined completely independently (and by different workers) and only the one using the absolute magnitude needs an assumption for the distance scale."
+ e Then using Eq. (, $\bullet $ Then using Eq. (
+15). we estimate the vertical velocity dispersion at one photometric scale leneth im the baud.,"15), we estimate the vertical velocity dispersion at one photometric scale length in the -band."
+ For this we need a value for the mass-to-light ratio aud we will discuss this first., For this we need a value for the mass-to-light ratio and we will discuss this first.
+ We found. that through. Eq. (," We found, that through Eq. ("
+13) Bottema’s relation (1) provides a value for Q(AL/£) of about 5.7.,13) Bottema's relation (1) provides a value for $Q (M/L)$ of about 5.7.
+ So we make a choice for Q rather than for AL/£., So we make a choice for $Q$ rather than for $M/L$.
+ It has become customary to asstume values of Q of order 2. mainly based onu the nuuerical simulations of Sellwood Carlbere (1981). who find thei disks to settle with Q~ 1.7 at all radii.," It has become customary to assume values of $Q$ of order 2, mainly based on the numerical simulations of Sellwood Carlberg (1984), who find their disks to settle with $Q \sim $ 1.7 at all radii."
+ In principle we can use the observed properties of the Galaxy to fix Q from Eq. (, In principle we can use the observed properties of the Galaxy to fix $Q$ from Eq. (
+17).,17).
+" We have (o,fon)?~0.5 Gn the solar uciglbourliood. but assuue for the sake of the argument also at R= 15) aud hh~ (see Sackett 1997 for a receut review). so that indeed Q~L7."," We have $(\sigma _{\rm z}/\sigma
+_{\rm R})^{2} \sim 0.5$ (in the solar neighbourhood, but assume for the sake of the argument also at $R = 1h$ ) and $h_{\rm z}/h \sim 0.1$ (see Sackett 1997 for a recent review), so that indeed $Q \sim 1.7$."
+ We will make the general assuiiptiou that Q = 2. In agreement with the considerations above: then (AL/L)p — 2m.," We will make the general assumption that $Q$ = 2, in agreement with the considerations above; then $(M/L)_{B}$ = 2.8."
+ The rotation velocity version of Dottemias empirical relations (Eq. (, The rotation velocity version of Bottema's empirical relations (Eq. (
+1)) cau provide further support for the choice of Q. aloug the lines of the discussion in van der Iruit Freeman (1986).,"1)) can provide further support for the choice of $Q$, along the lines of the discussion in van der Kruit Freeman (1986)."
+ In the first place we recall the condition for the prevention of swing amplification in disks (Toomre 1981). as reformulated by Sellwood (1983) where 5s is the number of spiral arius.," In the first place we recall the condition for the prevention of swing amplification in disks (Toomre 1981), as reformulated by Sellwood (1983) where $m$ is the number of spiral arms."
+ For a fat rotation curve this can be rewritten as With Eq. (, For a flat rotation curve this can be rewritten as With Eq. (
+1) this becomes Q>1.15i5..,1) this becomes $Q \gt 1.15 m$.
+ Cousideriueg -iat the coefficient in Eq. (, Considering that the coefficient in Eq. (
+1) has au uncertainty of order5%.. this tells us that we wave to assume (Q at least of order 2 to prevent strong barlike (:50—2) disturbances m the disk.,"1) has an uncertainty of order, this tells us that we have to assume $Q$ at least of order 2 to prevent strong barlike =2) disturbances in the disk."
+ A similar argument can be made using the egloha stability criterion of Efstathiou et al. (, A similar argument can be made using the global stability criterion of Efstathiou et al. (
+1982).,1982).
+ This criterion states that for a galaxy with a flat rotation curve and au expoucntial disk. global stability requires a dark halo ac Tere Aquas is the total mass of the disk.," This criterion states that for a galaxy with a flat rotation curve and an exponential disk, global stability requires a dark halo and Here $M_{\rm disk}$ is the total mass of the disk."
+ Tlis cau he rewritten as and. when evaluated at R=Lh. vields with Eq. (," This can be rewritten as and, when evaluated at $R = 1h$, yields with Eq. ("
+1) 0.69.702LAL. an therefore also implies that Q should be at least about 2.,"1) $0.69 \sqrt{Q} 
+\simgt 1.1$, and therefore also implies that $Q$ should be at least about 2."
+ Efstathiou et al., Efstathiou et al.
+ have also come to this conclusion for our Galaxy; with the use of local parameters for the solar ucighbourhood.," have also come to this conclusion for our Galaxy, with the use of local parameters for the solar neighbourhood."
+ Ihwiug adopted a value for Q and throueh this a value for (CM/L)p. we will have to convert it to (ML)j.," Having adopted a value for $Q$ and through this a value for $(M/L)_{B}$, we will have to convert it to $(M/L)_{I}$."
+ For this we need a (8B...£) colour for the disks., For this we need a $(B - I)$ colour for the disks.
+ From the fits of de Cais (1998) we find that the totaldisk magnitudes show a rather laree variation iu colour: for the sample used here (B.£) has a mean value of 1.9. but the ranis.," From the fits of de Grijs (1998) we find that the total magnitudes show a rather large variation in colour; for the sample used here $(B - I)$ has a mean value of 1.9, but the r.m.s."
+ scatter is O.S magnitudes., scatter is 0.8 magnitudes.
+ Iu. his discussion.de Cuijs (1998) suspects a systematic effect of the internal dust iu the disks (particularly ou the P-maguitudes. which is another reason for us to use the Vigg-version of the Bottema relation (Eq. (," In his discussion,de Grijs (1998) suspects a systematic effect of the internal dust in the disks (particularly on the -magnitudes, which is another reason for us to use the $V_{\rm rot}$ -version of the Bottema relation (Eq. ("
+1)) in our derivation iu the previous section).,1)) in our derivation in the previous section).
+ Iustead we turn to the discussion of de Jong (19961). who compares his surface photometry of less inclined spirals to star formation models.," Instead we turn to the discussion of de Jong (1996b), who compares his surface photometry of less inclined spirals to star formation models."
+ Frou his Table 3. we infer that for single burst models with solar metallicity and ages of 12 Car CM/L)p=2(U/L).," From his Table 3, we infer that for single burst models with solar metallicity and ages of 12 Gyr $(M/L)_{B} = 2
+(M/L)_{I}$."
+" So we will use an (AL/£), of 1.1.", So we will use an $(M/L)_{I}$ of 1.4.
+ There is a further refinement required., There is a further refinement required.
+ In order to take into account the fact that in late-tvpe galaxies the eas coutributes significautly to the gravitational force. we have to correct for a galaxw’s gas conteut as a function of IIubble type.," In order to take into account the fact that in late-type galaxies the gas contributes significantly to the gravitational force, we have to correct for a galaxy's gas content as a function of Hubble type."
+ Iu the following. we will discuss the observational data regarding the aud the Ty separately.," In the following, we will discuss the observational data regarding the and the $_2$ separately."
+ For 25 ofde Caijs’ sample galaxies observations are available. so hat we can estimate the eas-to-total disk mass.," For 25 of de Grijs' sample galaxies observations are available, so that we can estimate the gas-to-total disk mass."
+ For this we apply a correction of a factor 1/3 to the in order to take account of helimn aud use de Ciijs (1998) Fbaud photometry aud our adopted Ευ ratio of 2.8 (see below) to estimate the total disk mass., For this we apply a correction of a factor 4/3 to the in order to take account of helium and use de Grijs' (1998) -band photometry and our adopted $M/L_B$ ratio of 2.8 (see below) to estimate the total disk mass.
+" As a function of Hubble type we then find We find no dependence on rotation velocity: So. the mass is abot half the stellar mass in disks of Sls aud about similar t that iu σος and σας,"," As a function of Hubble type we then find We find no dependence on rotation velocity: So, the mass is about half the stellar mass in disks of Sb's and about similar to that in Sc's and Sd's."
+ But there is no depenudenuce on rotation velocity., But there is no dependence on rotation velocity.
+" But this is not what we need: we should use surface densities rather than disk. masses,", But this is not what we need; we should use surface densities rather than disk masses.
+ Now. the is usually more exteuded than the," Now, the is usually more extended than the"
+There are nevertheless some sections of spiral arms that appear to cross between two minima.,There are nevertheless some sections of spiral arms that appear to cross between two minima.
+"The stellar spiral arms associated with the N-body simulation are transient, so the minimum will eventually disappear.","The stellar spiral arms associated with the N-body simulation are transient, so the minimum will eventually disappear."
+" However the gaseous spiral arms still retain their spiral structure, even though the minimum dissolves."," However the gaseous spiral arms still retain their spiral structure, even though the minimum dissolves."
+ Eventually the gas reaches another minimum and joins a new spiral arm., Eventually the gas reaches another minimum and joins a new spiral arm.
+ The final panel in Fig., The final panel in Fig.
+ 3 shows a different section of spiral arm at a much later time frame., 3 shows a different section of spiral arm at a much later time frame.
+ The selected particles have already passed 3 spiral arms at this point., The selected particles have already passed 3 spiral arms at this point.
+" Gas particles highlighted with y> —2kpcare largely coincident with a potential minimum, and also corresponds to one of the dense arms of gas in Fig."," Gas particles highlighted with $y>-2$ kpc are largely coincident with a potential minimum, and also corresponds to one of the dense arms of gas in Fig."
+ 4., 4.
+" On the other hand, gas particles with «>—2 kpc lie between two minima or are forming the new spiral arm seen in the lower part of the figure."," On the other hand, gas particles with $x>-2$ kpc lie between two minima or are forming the new spiral arm seen in the lower part of the figure."
+ In Fig., In Fig.
+" 4, spiral arms lying between the potential minima are most evident when cold gas is present and the densities of the gas are higher."," 4, spiral arms lying between the potential minima are most evident when cold gas is present and the densities of the gas are higher."
+" A further consequence of gas retaining the spiral structure imposed by the potential is that at later times in the simulations (e.g. 3rd time frame, Fig."," A further consequence of gas retaining the spiral structure imposed by the potential is that at later times in the simulations (e.g. 3rd time frame, Fig."
+" 1), most of the gas is in the spiral arms."," 1), most of the gas is in the spiral arms."
+" Regions between the arms are relatively empty, and since little gas is entering the potential, the spiral arms correspond merely to dense regions of gas rather than actual spiral shocks."," Regions between the arms are relatively empty, and since little gas is entering the potential, the spiral arms correspond merely to dense regions of gas rather than actual spiral shocks."
+" In addition to the density, the velocity of the gas which is in the spiral arms largely follows that of the stellar component of the disc, with different pattern speeds emerging for different sections of spiral arm."," In addition to the density, the velocity of the gas which is in the spiral arms largely follows that of the stellar component of the disc, with different pattern speeds emerging for different sections of spiral arm."
+ The rotational velocities of the gas thus indicate that there is no fixed pattern speed or co-rotation radius for the spiral perturbation., The rotational velocities of the gas thus indicate that there is no fixed pattern speed or co-rotation radius for the spiral perturbation.
+"the supergiant ellipticals that lie at the centers of rich clusters of galaxies, which are the very objects that are the main landmarks in the Hubble flow.","the supergiant ellipticals that lie at the centers of rich clusters of galaxies, which are the very objects that are the main landmarks in the Hubble flow."
+ In this paper we present our new measurement of the GCLF turnover for the globular cluster system in NGC 4874. the central cD elliptical in the Coma cluster. and use it to estimate Hy.," In this paper we present our new measurement of the GCLF turnover for the globular cluster system in NGC 4874, the central cD elliptical in the Coma cluster, and use it to estimate $H_0$."
+" The bright end of the GCLF in NGC 4874 has already been studied with ground-based imaging, which indicated that it indeed has a large globular cluster system (Harris 1987:; Thompson&Valdes 1987: Blakeslee&Tonry 1995)."," The bright end of the GCLF in NGC 4874 has already been studied with ground-based imaging, which indicated that it indeed has a large globular cluster system \cite{har87}; ; \cite{tho87}; \cite{bla95}) )."
+" With the much deeper HST photometric limits, we could therefore confidently expect to garner a huge population of clusters to define the GCLE."," With the much deeper HST photometric limits, we could therefore confidently expect to garner a huge population of clusters to define the GCLF."
+ Our raw dataset consists of 18 V (F606W) and 10 7 (F814W) images of various exposure lengths (see Table 1). taken 1997 August 16 and 24 (program GO-5905) with the WFPC2 camera.," Our raw dataset consists of 18 $V$ (F606W) and 10 $I$ (F814W) images of various exposure lengths (see Table \ref{tab:obs}) ), taken 1997 August 16 and 24 (program GO-5905) with the WFPC2 camera."
+The long exposures were sub-pixel-shifted (dithered) in a pentagonal pattern by fractional pixel amounts in order to reconstruct clean composite images free from cosmic-ray contamination and bad-pixel artifacts.,The long exposures were sub-pixel-shifted (dithered) in a pentagonal pattern by fractional pixel amounts in order to reconstruct clean composite images free from cosmic-ray contamination and bad-pixel artifacts.
+" The V exposures, totaling 20940 sec, were taken to probe the GCLF deeply enough to resolve the turnover point. while the 7 exposures, totaling 8720 sec, were used to define the color (metallicity) distribution for the brighter end of the cluster system."," The $V$ exposures, totaling 20940 sec, were taken to probe the GCLF deeply enough to resolve the turnover point, while the $I$ exposures, totaling 8720 sec, were used to define the color (metallicity) distribution for the brighter end of the cluster system."
+" The color distribution, spatial structure of the GCS, and the specific frequency will be discussed in Paper Η (Harrisetal. 1999))."," The color distribution, spatial structure of the GCS, and the specific frequency will be discussed in Paper II \cite{har99a}) )."
+ Here. we analyze the GCLF and use it to estimate the distance to Coma and thus Hi.," Here, we analyze the GCLF and use it to estimate the distance to Coma and thus $H_0$."
+" To maximize the total elobular cluster population falling within the WFPC? field of view, we placed the center of the PCI CCD on the nucleus of NGC 4874."," To maximize the total globular cluster population falling within the WFPC2 field of view, we placed the center of the PC1 CCD on the nucleus of NGC 4874."
+" A few other large elliptical galaxies projected near the center of Coma fall on the outskirts of the WF2.3.4 CCDs: however, these proved not to have significant numbers of globular clusters of their own and thus did not contaminate the NGC 4874 sample."," A few other large elliptical galaxies projected near the center of Coma fall on the outskirts of the WF2,3,4 CCDs; however, these proved not to have significant numbers of globular clusters of their own and thus did not contaminate the NGC 4874 sample."
+" Small areas surrounding them were, in any case, masked out in all subsequent data analysis."," Small areas surrounding them were, in any case, masked out in all subsequent data analysis."
+ We first retrieved the raw data from the HST archive located at theCADC?., We first retrieved the raw data from the HST archive located at the.
+". The CADC pipeline preprocessed the images at this point, with the best calibration images then available."," The CADC pipeline preprocessed the images at this point, with the best calibration images then available."
+ We then combined the exposures in pairs (the pairs of long exposures within each orbit) to define a first set of frames reasonably tree of cosmic-ray contamination., We then combined the exposures in pairs (the pairs of long exposures within each orbit) to define a first set of frames reasonably free of cosmic-ray contamination.
+ The vast majority of detected objects on the frames are the globular clusters in the halo of NGC 4874., The vast majority of detected objects on the frames are the globular clusters in the halo of NGC 4874.
+" At the ~100—Mpc distance of Coma they appear as unresolved point sources even in the PCI frames, so it is readily possible to perform conventional point-spread function (PSP) photometry on the frames."," At the $\sim 100-$ Mpc distance of Coma they appear as unresolved point sources even in the PC1 frames, so it is readily possible to perform conventional point-spread function (PSF) photometry on the frames."
+ We constructed an independent PSF for cach image and each of the four WFPC2 CCDs., We constructed an independent PSF for each image and each of the four WFPC2 CCDs.
+" With DAOPHOT and ALLSTAR (Stetson1994)) we then generated separate lists of candidate starlike (that is, unresolved) objects on cach frame."," With DAOPHOT and ALLSTAR \cite{ste94}) ) we then generated separate lists of candidate starlike (that is, unresolved) objects on each frame."
+ These coordinate lists were used to determine the (small) offsets and rotations to map the images onto a common re-registered coordinate system., These coordinate lists were used to determine the (small) offsets and rotations to map the images onto a common re-registered coordinate system.
+" Next, Stetson’s (1994) MONTAGE? code was used to define a “master” image in each filter as the median (i... 50th percentile) of the individual exposures."," Next, Stetson's (1994) MONTAGE2 code was used to define a “master” image in each filter as the median (i.e., 50th percentile) of the individual exposures."
+" Finally, we summed the master V and / images to generate a single deep. contamination-tree image which maximized the available flux for object detection."," Finally, we summed the master $V$ and $I$ images to generate a single deep, contamination-free image which maximized the available flux for object detection."
+" The photometry must also be designed to avoid false detections and nonstellar objects (noise spikes, small, faint background galaxies, or even compact dwarf galaxies within Coma itself)."," The photometry must also be designed to avoid false detections and nonstellar objects (noise spikes, small, faint background galaxies, or even compact dwarf galaxies within Coma itself)."
+" The real and artificial-star measurements were therefore combined with image shape parameters (the DAOPHOT parameters SHARP and 4, and the radial image moment 7, defined by Kron1980. and Harrisetal. 1991))."," The real and artificial-star measurements were therefore combined with image shape parameters (the DAOPHOT parameters SHARP and $\chi$, and the radial image moment $r_1$ defined by \cite{kro80} and \cite{har91}) )."
+" Several numerical trials were carried out with the artificial-star data to vary the detection threshold (in units of o,. the RMS scatter of the sky background) and the shape selection parameters, and thus to determine the highest values of these thresholds which would not deteriorate the detection efficiency of genuine starlike objects (see Figure 1)."," Several numerical trials were carried out with the artificial-star data to vary the detection threshold (in units of $\sigma_s$, the RMS scatter of the sky background) and the shape selection parameters, and thus to determine the highest values of these thresholds which would not deteriorate the detection efficiency of genuine starlike objects (see Figure \ref{fig:image_moments}) )."
+" We adopted a DAOPHOT/FIND detection threshold of 3.50, above sky (c.f. Stetson 1987)."," We adopted a DAOPHOT/FIND detection threshold of $3.5 \sigma_s$ above sky (c.f. \cite{ste87}) ),"
+ a level which recovered virtually all the brighter input artificial stars while adding almost no false detections from noise., a level which recovered virtually all the brighter input artificial stars while adding almost no false detections from noise.
+" The adopted boundaries for V. SHARP, and r, resulted in the culling of about of the artificial stars, mostly at the faintest levels where the distinction between starlike and nonstellar objects by image shape classification also becomes difficult."," The adopted boundaries for $\chi$, SHARP, and $r_1$ resulted in the culling of about of the artificial stars, mostly at the faintest levels where the distinction between starlike and nonstellar objects by image shape classification also becomes difficult."
+" At the faint end as well. some unwanted nonstellar objects end up being scattered into the “starlike” category, but these must be statistically removed from the final GCLF by subtraction of the background luminosity function (see below)."," At the faint end as well, some unwanted nonstellar objects end up being scattered into the “starlike” category, but these must be statistically removed from the final GCLF by subtraction of the background luminosity function (see below)."
+" With the master list of starlike objects now determined, we used the ALLFRAME code (Stetson 1994)) to measure these objects on the original set of F6060Wand F814W images. employing the individual PSFs for each frame as determined"," With the master list of starlike objects now determined, we used the ALLFRAME code \cite{ste94}) ) to measure these objects on the original set of F606Wand F814W images, employing the individual PSFs for each frame as determined"
+"How galaxies formand the stas within hemd8 one of the major open questions of uoderu costuology,",How galaxies form—and the stars within them—is one of the major open questions of modern cosmology.
+ Early-type galaxies (ETC) iost a large fraction of the stellar mass iu todays universe. ando typically show no evidence for ongoing star∙ formation. so∙↴∖↴↑↸∖↕⋜∐⋅⋜↧∶↴∙⊾↸∖↴∖↴⋜⋯≺⋜∏⋝∏∐≼⋜⋯↸⊳↸∖↴∖↴↕∐⊏⊺≼∶↴∖↴∙⊺∪≼⋜↧∙↖↽ they are natural here arects for the investigation of how and when galaxy asscliubly aud star formation occurred in ho past.," Early-type galaxies (ETGs) host a large fraction of the stellar mass in today's universe, and typically show no evidence for major ongoing star formation, so they are natural targets for the investigation of how and when galaxy assembly and star formation occurred in the past."
+" The task is a difficult one, which has constuned a large amount of observational aud heoretical∙ effort ∙∙in the past several decades."," The task is a difficult one, which has consumed a large amount of observational and theoretical effort in the past several decades."
+ All hat work cannot be fairly sununarized iu a such short |article., All that work cannot be fairly summarized in a such short article.
+ This review is thus narrowly‘ focused ou What has been learned frou stellar ages and based on iuteerated ποτ. studies.. .roni and the new questions. raised. by recent evideuce.," This review is thus narrowly focused on what has been learned from stellar ages and abundances based on integrated light studies, and the new questions raised by recent evidence."
+. Apologies]© ©eo to manytm hard-workingoct© colleaguesUo for auv muportaut onmussions., Apologies go to many hard-working colleagues for any important omissions.
+ Tistorically. the debate on the star forination hnistorv of ETCs was framed in terms of two colupeting scenarios: hierarchical clustering (6.8. White Rees 1978. Searle Zinn 1978). according to which these galaxies were assembled through the mereiue of less massive structures formed at vel redshift: aud monolithic dissipative collapse (c.g.. Egeeu 1962. Larson. 1971). whereby uassive ETGs were formed at very high redshift jv measof a rapid eravitational collapse.," Historically, the debate on the star formation history of ETGs was framed in terms of two competing scenarios: hierarchical clustering (e.g., White Rees 1978, Searle Zinn 1978), according to which these galaxies were assembled through the merging of less massive structures formed at high redshift; and monolithic dissipative collapse (e.g., Eggen 1962, Larson 1974), whereby massive ETGs were formed at very high redshift by means of a rapid gravitational collapse."
+ Deciding )etween these two scenarios was one of the uai imnofivatious behiud attempts to iueasmre ua, Deciding between these two scenarios was one of the main motivations behind attempts to measure stellar ages and abundances in ETGs.
+jor is little question that galaxies formed üerarchicalle in a A-CDAL universe. and while hat historical debate has been settled. studies of unresolved stellar populations have acquired renewed portance. as they provide umch needed constraints to iucreasinely sophisticated galaxy ‘ormation models.," Today there is little question that galaxies formed hierarchically in a $\Lambda$ -CDM universe, and while that historical debate has been settled, studies of unresolved stellar populations have acquired renewed importance, as they provide much needed constraints to increasingly sophisticated galaxy formation models."
+ Early attempts at dating. stellar populations. .applications of⋅⋅ stellar⋅ population svuthesis. abundances. uodels to observations. of .inteerated light were based on photometricor low-resolution spectroplotonietric observatious. (6.5e. O'Connell 1980. Comm 1981. Renzini 1986) aud hielh-resolution photographic spectroscopy (e.g... Rose 1985).," Early attempts at dating stellar populations from applications of stellar population synthesis models to observations of integrated light were based on photometric or low-resolution spectrophotometric observations (e.g., O'Connell 1980, Gunn 1981, Renzini Buzzoni 1986) and high-resolution photographic spectroscopy (e.g., Rose 1985)."
+ Thev led to pronüsing. vet not cutively conclusive results. due to limitations of the carly models and/or uncertainties associated with the ageauetallicitv degeneracy (0.g.. Reuzini 1986. Wortley 1991).," They led to promising, yet not entirely conclusive results, due to limitations of the early models and/or uncertainties associated with the age-metallicity degeneracy (e.g., Renzini 1986, Worthey 1994)."
+ The latter is a inandfestatiou of the similar dependence of the temperatures of nian sequence, The latter is a manifestation of the similar dependence of the temperatures of main sequence
+Regarding the Sy and control sample. we found à similar incidence of bars. rings. asymmetries. and close companions. considered both on an individual basis and all together. with the Sy bars somewhat weaker.,"Regarding the Sy and control sample, we found a similar incidence of bars, rings, asymmetries, and close companions, considered both on an individual basis and all together, with the Sy bars somewhat weaker."
+ Thus. our results imply that the fueling of Sy nuclei is not directly related to large-scale mechanisms operating over the bulk of the gas.," Thus, our results imply that the fueling of Sy nuclei is not directly related to large-scale mechanisms operating over the bulk of the gas."
+ There are some hints. however. of a link between them.," There are some hints, however, of a link between them."
+ First. it is generally believed that the required fuel is a tiny fraction of the gas in the inner few ppc. especially of spiral galaxies. and angular momentum reduction is the major challenge (e.g..Jogee2006).," First, it is generally believed that the required fuel is a tiny fraction of the gas in the inner few pc, especially of spiral galaxies, and angular momentum reduction is the major challenge \citep[e.g.,][]{J_06}."
+. For instance. typical molecular gas mass of =10°M was reported for the central regions of most galaxies (e.g..Garefa-Burilloetal.2005).," For instance, typical molecular gas mass of $\approx\,$$10^8\,M_\odot$ was reported for the central regions of most galaxies \citep[e.g.,][]{GCS04_05}."
+ A part of this gas is expected to have resulted from secular evolution., A part of this gas is expected to have resulted from secular evolution.
+ A higher molecular gas concentration was found in the central kiloparsee of barred galaxies than in non-barred ones (Sakamotoetal.1999;Sheth2005.seealsoReganetal. 200601: according to the first authors. more than half of the central gas was driven there by the bar.," A higher molecular gas concentration was found in the central kiloparsec of barred galaxies than in non-barred ones \citep[][see also Regan et~al. 2006]{SOI_99,SVR_05}; according to the first authors, more than half of the central gas was driven there by the bar."
+ The gas in nuclear rings. the most evident tracers of recent gas inflow. can be brought under the influence of the SMBH by viscous torques in the scenario of Garefa-Burilloetal.(2005).," The gas in nuclear rings, the most evident tracers of recent gas inflow, can be brought under the influence of the SMBH by viscous torques in the scenario of \citet[][]{GCS04_05}."
+. Furthermore. higher central gas concentration has been associated with interactions and mergers (e.g..Georgakakisetal.2000:Smith2007).," Furthermore, higher central gas concentration has been associated with interactions and mergers \citep[e.g.,][]{GFN_00,SSH_07}."
+" Second. generally weaker bars in Sy than in inactive galaxies have been associated with larger amounts of cold gas in their host galaxies in the framework of central mass concentrations that could destroy x, bar orbits (Shlosmanetal. 2000)."," Second, generally weaker bars in Sy than in inactive galaxies have been associated with larger amounts of cold gas in their host galaxies in the framework of central mass concentrations that could destroy $x_1$ bar orbits \citep[][]{SPK_00}."
+. It has been shown. however. that bars are less fragile than previously thought. and the mass of the central concentration required to dissolve the bar must be very high (e.g..ΓιαShen&Sellwood2004:Debattistaetal.2006;Marinova&Jogee 2007).," It has been shown, however, that bars are less fragile than previously thought, and the mass of the central concentration required to dissolve the bar must be very high \citep[e.g.,][]{SS_04,DMC_06,MJ_07}."
+. Alternatively. the main destruction mechanism could be the transfer of angular momentum from the gas inflow to the bar (e.g..Bournaudetal.2005).. especially in the presence of radiative cooling (e.g..Debattistaetal.2006).," Alternatively, the main destruction mechanism could be the transfer of angular momentum from the gas inflow to the bar \citep[e.g.,][]{BCS_05}, especially in the presence of radiative cooling \citep[e.g.,][]{DMC_06}."
+. Thus. the weaker Sy bars may be related to the generally larger cold gas amounts reported in their disks (e.g..Huntetal.1999.seealsoHoetal.2008) in the context of angular momentum transfer.," Thus, the weaker Sy bars may be related to the generally larger cold gas amounts reported in their disks \citep[e.g.,][see also Ho et~al. 2008]{HMM2_99} in the context of angular momentum transfer."
+ The relatively low accretion rates of Sy nuclei prompt a variety of small-scale processes able to drive the circumnuclear gas down to the very centre (e.g..Martini2004).," The relatively low accretion rates of Sy nuclei prompt a variety of small-scale processes able to drive the circumnuclear gas down to the very centre \citep[e.g.,][]{M_04}."
+. This could be the main reason for a lack of a universal morphological pattern on these scales (e.g..Garefa-Burilloetal.2004).," This could be the main reason for a lack of a universal morphological pattern on these scales \citep[e.g.,][]{GCS_04}."
+.. Seyfert activity. however. has been associated with the presence of dust (SimóesLopesetal.2007) and more disturbed gaseous kinematics (Dumasetal.2007) in the circumnuclear regions.," Seyfert activity, however, has been associated with the presence of dust \citep[][]{SSF_07} and more disturbed gaseous kinematics \citep[][]{DME_07} in the circumnuclear regions."
+ In this regard. we started a study of the circumnuclear regions of a sample of Sy galaxies using HST archival images.," In this regard, we started a study of the circumnuclear regions of a sample of Sy galaxies using HST archival images."
+ The circumnuclear structures of 3352 and 5590 are the first results of this research., The circumnuclear structures of 352 and 590 are the first results of this research.
+ In the framework of our results we have started a study of the circumnuclear regions of à sample of Sy galaxies using HST archival images., In the framework of our results we have started a study of the circumnuclear regions of a sample of Sy galaxies using HST archival images.
+ As first results of this research. we revealed a nuclear bar and ring in 3352 and nuclear dust lanes in 5590.," As first results of this research, we revealed a nuclear bar and ring in 352 and nuclear dust lanes in 590."
+limb darkening and gravity darkening tend to limit their contribution to the line profiles and affect the reconstruction accuracy of low-latitude magnetic features.,limb darkening and gravity darkening tend to limit their contribution to the line profiles and affect the reconstruction accuracy of low-latitude magnetic features.
+ Synthetic Stokes V profiles are computed for all observed rotation phases and compared to the observations., Synthetic Stokes V profiles are computed for all observed rotation phases and compared to the observations.
+ The model adjustment is iterative and based on a maximum entropy algorithm (Skilling Bryan 1984)., The model adjustment is iterative and based on a maximum entropy algorithm (Skilling Bryan 1984).
+ The version of the code used here makes a projection of the surface magnetic field onto a spherical harmonics frame (Donati et al., The version of the code used here makes a projection of the surface magnetic field onto a spherical harmonics frame (Donati et al.
+" 2006), with the magnetic field geometry resolved into its poloidal and toroidal component (Chandrasekhar 1961)."," 2006), with the magnetic field geometry resolved into its poloidal and toroidal component (Chandrasekhar 1961)."
+" We limit the spherical harmonics expansion to €«10, since no improvement in the fit to the data is achieved by increasing further the maximum allowed value for ¢."," We limit the spherical harmonics expansion to $\ell<10$, since no improvement in the fit to the data is achieved by increasing further the maximum allowed value for $\ell$."
+ The first step to reconstruct a relevant topology of the surface field consists in determining the stellar rotation period., The first step to reconstruct a relevant topology of the surface field consists in determining the stellar rotation period.
+" To do so, we follow the approach of Petit et al. ("," To do so, we follow the approach of Petit et al. ("
+"2002) where a set of magnetic maps is calculated, assuming for each map a different value for the rotation period.","2002) where a set of magnetic maps is calculated, assuming for each map a different value for the rotation period."
+ We impose a constant entropy for all images and calculate a reduced hhereafter) by comparing the set of synthetic Stokes V profiles produced by ZDI to the observed time-series of profiles., We impose a constant entropy for all images and calculate a reduced hereafter) by comparing the set of synthetic Stokes V profiles produced by ZDI to the observed time-series of profiles.
+ The resulting vvariations (plotted in Fig. 3)), The resulting variations (plotted in Fig. \ref{fig:period}) )
+ are recorded over the range of rotation periods to determine the period value producing the best magnetic model (identified by the lowest value of the reduced y? goodness-of-fit parameter)., are recorded over the range of rotation periods to determine the period value producing the best magnetic model (identified by the lowest value of the reduced $\chi^2$ goodness-of-fit parameter).
+" Here, we scan 300 values of the period between 0.4 d and 1 d, a range that"," Here, we scan 300 values of the period between 0.4 d and 1 d, a range that"
+"was used to place upper limits on e for 28 isolated pulsars with f,>25 Lz (TheLIGOB.Abbottetal.2004a).",was used to place upper limits on $\epsilon$ for 28 isolated pulsars with $f_* > 25$ Hz \citep{lig04b}.
+". Our results indicate that these (time- and frequency-domain) search strategies must be revised to include the signal al f. (if the mountain is static) and even to collect signal within a bandwidth A/ centered al [, and 2f, Gf the mountain oscillates).", Our results indicate that these (time- and frequency-domain) search strategies must be revised to include the signal at $f_*$ (if the mountain is static) and even to collect signal within a bandwidth $\Delta f$ centered at $f_*$ and $2f_*$ (if the mountain oscillates).
+ This remains true under several of the evolutionary scenarios outlined above when precession is included. depending on the (unknown) competitive balance between driving and damping.," This remains true under several of the evolutionary scenarios outlined above when precession is included, depending on the (unknown) competitive balance between driving and damping."
+ The analvsis in (his paper clisreeards the [act that LIGO I] will be tunable., The analysis in this paper disregards the fact that LIGO II will be tunable.
+" It is important to redo the SNR caleulations with realistic (unable noise curves. (o investigate whether the likelihood of detection is maximized by observing near f, or 2/,."," It is important to redo the SNR calculations with realistic tunable noise curves, to investigate whether the likelihood of detection is maximized by observing near $f_*$ or $2f_*$."
+ We also do nol consider several physical processes that affect magnetic burial. such as sinking of accreted malerial Ohnmüe dissipation. or Ilall currents: their importance is estimated roughly by Alelatos&Payne (2005)..," We also do not consider several physical processes that affect magnetic burial, such as sinking of accreted material, Ohmic dissipation, or Hall currents; their importance is estimated roughly by \citet{mel05}. ."
+ Finally. Doppler shifts due to the Earth's orbit and. rotation (e.g.Donazzola&Gourgoulhon1996) are neglected. as are slow secular cdrifts in sensitivity during acoherent integration.," Finally, Doppler shifts due to the Earth's orbit and rotation \citep[e.g.][]{bon96} are neglected, as are slow secular drifts in sensitivity during acoherent integration."
+L>0.05L galaxies are hosts to absorbing gas halos characterized by a covering fraction of unity and a spherical geometry which truncates at R(L).,$L> 0.05L^{\ast}$ galaxies are hosts to absorbing gas halos characterized by a covering fraction of unity and a spherical geometry which truncates at $R(L)$.
+" Examination of this now ""standard model"" has been the subject of several theoretical studies (e.g..Charlton&Churchill1996:MoMiralda-Escude1996;Lin&Zou2001 )."," Examination of this now “standard model” has been the subject of several theoretical studies \citep[e.g.,][]{cc96,mo96,lin01}."
+. Guillemin&Bergeron(1997) determined a steeper value of ./20.28 for the B-band luminosity obtained from a best fit to the upper envelope of the distribution of impact parameters of 26 absorbing galaxies., \citet{gb97} determined a steeper value of $\beta = 0.28$ for the B–band luminosity obtained from a best fit to the upper envelope of the distribution of impact parameters of 26 absorbing galaxies.
+ They found Δ.=67 kpe., They found $R_{\ast} = 67$ kpc.
+ Using a reverse approach of establishing foreground galaxy redshifts and then searching for absorption in the spectra of background quasars yields results inconsistent with a covering fraction of unity., Using a reverse approach of establishing foreground galaxy redshifts and then searching for absorption in the spectra of background quasars yields results inconsistent with a covering fraction of unity.
+ For example. Bowenetal.(1995) identified 17 low-redshift galaxies with background quasar probing an impact parameter range between 3—162 kpe.," For example, \citet{bowen95} identified 17 low–redshift galaxies with background quasar probing an impact parameter range between $3-162$ kpc."
+ Galaxies that were probed at impact parameters greater than 13 kpe had no absorption in the halo (W.(2796)50.40—0.9 A). however. four of the six galaxies within 13 kpe of the halo produced absorption.," Galaxies that were probed at impact parameters greater than 13 kpc had no absorption in the halo $W_{r}(2796) \geq
+0.40-0.9$ ), however, four of the six galaxies within 13 kpc of the halo produced absorption."
+" For intermediate redshift galaxies. Bechtold&Ellingson(1992) reported a covering fraction f.~0.25 for W,(2796)>0.26 for eight galaxies with D<85 kpe."," For intermediate redshift galaxies, \citet{bechtold92} reported a covering fraction $f_c \simeq
+0.25$ for $W_{r}(2796) \geq 0.26$ for eight galaxies with $D
+\leq 85$ kpc."
+ Also. Tripp&Bowen(2005). reported f.0.5 for W(2796)>0.15 for ~20 galaxies with D<50 kpe.," Also, \citet{tripp-china} reported $f_c \sim 0.5$ for $W_{r}(2796) \geq 0.15$ for $\sim 20$ galaxies with $D \leq 50$ kpc."
+" These results are also consistent with the findings of Churchill.Kaeprzak.&Steidel(2005) who reported very weak absorption. W,(2796)<0.3A. well inside the R(L) boundary of bright galaxies: these galaxies would be classified as ""non-absorbers"" in previous surveys."," These results are also consistent with the findings of \citet{cwc-china} who reported very weak absorption, $W_{r}(2796) < 0.3$, well inside the $R(L)$ boundary of bright galaxies; these galaxies would be classified as “non–absorbers” in previous surveys."
+ They also report W.(2796)>| absorption out to ~2R(L)., They also report $W_r(2796) > 1$ absorption out to $\simeq 2 R(L)$.
+ All these results suggest that there are departures from the standard model. that the covering fraction of absorbing gas is less than unity. and that the halo sizes and the distribution of the gas appear to diverge from the R(L) relation with spherical geometry.," All these results suggest that there are departures from the standard model, that the covering fraction of absorbing gas is less than unity, and that the halo sizes and the distribution of the gas appear to diverge from the $R(L)$ relation with spherical geometry."
+ Another approach to understanding halo sizes and gas distributions is to determine the. statistical properties of absorbing gas and then compute the statistical cross section from the redshift path density. dN/dz (seeLanzettaetal. 1995)..," Another approach to understanding halo sizes and gas distributions is to determine the statistical properties of absorbing gas and then compute the statistical cross section from the redshift path density, $dN/dz$ \citep[see][]{lanzetta95}. ."
+ The downfall of this method is that a galaxy lummosity function must be adopted in order to estimate R..., The downfall of this method is that a galaxy luminosity function must be adopted in order to estimate $R_{\ast}$.
+" Nestoretal.(2005) acquired a sample of over 1300 absorption systems. with W,(2796)150.3 from the Sloan Digital Sky Survey (SDSS)."," \citet{nestor05} acquired a sample of over 1300 absorption systems, with $W_{r}(2796) \geq 0.3$ from the Sloan Digital Sky Survey (SDSS)."
+" Using the K-band Holmbere-like luminosity scaling and luminosity function of MUNICS (Droryetal.2003).. Nestor computed A.=60—100 kpe for adopted minimum luminosity cutoffs of L,,5,=0.001— 0.253L""."," Using the $K$ –band Holmberg--like luminosity scaling and luminosity function of MUNICS \citep{drory03}, Nestor computed $R_{\ast} = 60-100$ kpc for adopted minimum luminosity cutoffs of $L_{min}=0.001-0.25L^{\ast}$ ."
+ They found no redshift evolution of R. over the explored range of 0.3«z1.2., They found no redshift evolution of $R_{\ast}$ over the explored range of $0.3\leq z \leq 1.2$.
+ Zibettietal.(2007) studied the statistical photometricproperties of ~2800 absorbers in quasar fields imaged with SDSS., \citet{zibetti06} studied the statistical photometricproperties of $\sim2800$ absorbers in quasar fields imaged with SDSS.
+ Using the method of image stacking. they detected low-level surface brightness (SB) azimuthally about the quasar.," Using the method of image stacking, they detected low–level surface brightness (SB) azimuthally about the quasar."
+ The SB profiles follow a decreasing power law with projected distance away from the quasar out to 100—200 kpe., The SB profiles follow a decreasing power law with projected distance away from the quasar out to $100-200$ kpc.
+ These results imply that absorption selected galaxies may reside out to projected distances of 200 kpe., These results imply that absorption selected galaxies may reside out to projected distances of 200 kpc.
+ However. it is worth noting that the extended light profiles may be an artifact of clustering of galaxies.," However, it is worth noting that the extended light profiles may be an artifact of clustering of galaxies."
+ Cluster companions of the absorbing galaxies could extend the observed light profile over hundreds of stacked images., Cluster companions of the absorbing galaxies could extend the observed light profile over hundreds of stacked images.
+ Thus. one would infer that absorbing galaxies are present at a larger impact parameters than would be found in direct observation of individual galaxies.," Thus, one would infer that absorbing galaxies are present at a larger impact parameters than would be found in direct observation of individual galaxies."
+ Motivated by recent expectations from simulations that halo gas is dynamically complex and sensitive to the physics of galaxy formation. we investigate the standard halo model of absorbers.," Motivated by recent expectations from simulations that halo gas is dynamically complex and sensitive to the physics of galaxy formation, we investigate the standard halo model of absorbers."
+ We also aim to provide updated constraints on f. and ./ for galaxy formation simulations., We also aim to provide updated constraints on $f_c$ and $\beta$ for galaxy formation simulations.
+ In. this paper. we demonstrate that f£.<I and question the validity of the Holmbere-like luminosity scaling (Eq. 1)).," In this paper, we demonstrate that $f_c < 1$ and question the validity of the Holmberg–like luminosity scaling (Eq. \ref{eq:rl}) )."
+" Using high resolution quasar spectra. we explore absorption strengths to an order of magnitude more sensitive than previous surveys which allow us to re-identify non-absorbing galaxies as ""weak"" absorbing galaxies."," Using high resolution quasar spectra, we explore absorption strengths to an order of magnitude more sensitive than previous surveys which allow us to re–identify non–absorbing galaxies as “weak” absorbing galaxies."
+ In 2 we describe our sample and analysis., In \ref{sec:data} we describe our sample and analysis.
+ In 3.. we present new calculations of the statistical absorber radius computed using the statistically measured absorption path density dN/dz and the Schechter luminosity function.," In \ref{sec:results}, we present new calculations of the statistical absorber radius computed using the statistically measured absorption path density $dN/dz$ and the Schechter luminosity function."
+ We then compare these values to the empirical results of S95 and to a sample of known absorption selected galaxies with measured luminosities and impact parameters., We then compare these values to the empirical results of S95 and to a sample of known absorption selected galaxies with measured luminosities and impact parameters.
+ We also examine how individual halos behave with respect to the statistical halo., We also examine how individual halos behave with respect to the statistical halo.
+ In 4.. we discuss the properties and distribution of gas in halos.," In \ref{sec:dis}, we discuss the properties and distribution of gas in halos."
+ Our concluding remarks are in 5.., Our concluding remarks are in \ref{sec:conclusion}.
+ We have constructed a sample of 37 galaxies (0.3«=<1.0). with spectroscopically confirmed redshifts. selected by the presence of absorption in quasar spectra.," We have constructed a sample of 37 galaxies $0.3<z<1.0$ ), with spectroscopically confirmed redshifts, selected by the presence of absorption in quasar spectra."
+ The absorption properties were measured from HIRES/Keck (Vogtetal.1994). and UVES/VLT (Dekkeretal.2000) spectra.," The absorption properties were measured from HIRES/Keck \citep{vogt94} and UVES/VLT \citep{dekker00}
+ spectra."
+" The À2796 profileshave been presented in Churchill.Kaeprzak.&Steidel(2005).. where the detection limit is W,(2796)>0.02 (5 c)."," The $\lambda 2796$ profileshave been presented in \citet{cwc-china}, where the detection limit is $W_r(2796)\geq
+0.02$ (5 $\sigma$ )."
+ Galaxy properties were measured from F702W or F814WWFEPC-2/HST images of the quasar fields., Galaxy properties were measured from F702W or F814W images of the quasar fields.
+ Images of the galaxies. along with further details of the sample selection. data. and data analysis. can be found in Kaeprzaketal.(2007a).," Images of the galaxies, along with further details of the sample selection, data, and data analysis, can be found in \citet{kacprzak07}."
+. Galaxy absolute magnitudes. Mj. were determined from the A-corrected observed nose Or Aipsigiy adopted from Kaeprzaketal.(2007a).," Galaxy absolute magnitudes, $M_B$, were determined from the $k$ –corrected observed $m_{F702W}$ or $m_{F814W}$ adopted from \citet{kacprzak07}."
+. The A-corrections were computed using the formalism of Kim.Goobar.&Perlmutter(1996) based upon the spectral energy distribution (SED) templates of Kinneyetal.(1996)., The $k$ –corrections were computed using the formalism of \citet{kim96} based upon the spectral energy distribution (SED) templates of \citet{kinney96}.
+. The adopted SED for each galaxy was based upon its rest-frame B—K color (SDP94)., The adopted SED for each galaxy was based upon its rest–frame $B-K$ color (SDP94).
+ For galaxies with no color information. we adopteda Sb SED which is consistent with average color of absorbing galaxies (SDP94:Zibettietal.2007).," For galaxies with no color information, we adopteda Sb SED which is consistent with average color of absorbing galaxies \citep[SDP94;][]{zibetti06}."
+. Our K-corrections are consistent with those from the literature (Kim.Goobar.&Perlmut-ter1996:Fukugita.Shimasaku.," Our $k$ –corrections are consistent with those from the literature \citep{kim96,fukugita95}."
+"&Ichikawa 1995).. B-band luminosities were computed using the DEEP? optimal Mj, of Faberetal.(2007.Table2) in the redshift bin appropriate for each galaxy.", $B$ –band luminosities were computed using the DEEP2 optimal $M^{\ast}_B$ of \citet[][Table~2]{faber05} in the redshift bin appropriate for each galaxy.
+ Mj ranges from —21.07 (£z;2 0.3) to 21.54 (is) =1.1)., $M^{\ast}_B$ ranges from $-21.07$ $\left< z \right> =0.3$ ) to $-21.54$ $\left< z\right> =1.1$ ).
+ We compute the halo gascross section determined from the redshift path density. whereΑς is the statistical absorber radiusfor an L galaxy. and d is the number density of galaxies.," We compute the halo gascross section determined from the redshift path density, where$R_{\rm x}$ is the statistical absorber radiusfor an $L^{\ast}$ galaxy, and $\Phi^{\ast}$ is the number density of $L^{\ast}$ galaxies."
+ RI= f.R7. where R. is the covering fraction correctedL absorbing halo radius.," $R^2_{\rm x}=f_c R^2_{\ast}$ , where $R_{\ast}$ is the covering fraction corrected absorbing halo radius."
+ Note here that we make a distinction between Αν. which is derived from the redshift path density. and R... which is a physical cross section of the absorbing gas accounting for the," Note here that we make a distinction between $R_{\rm x}$ , which is derived from the redshift path density, and $R_{\ast}$ , which is a physical cross section of the absorbing gas accounting for the"
+any process that uses a sieuificance-based detection threshold.,any process that uses a significance-based detection threshold.
+ Were. we briefly set out a ecισα. recΠρο to use in more complicated cases.," Here, we briefly set out a general recipe to use in more complicated cases."
+ For complex detection algoritlius. sole of the steps may require Moute Carlo methods.," For complex detection algorithms, some of the steps may require Monte Carlo methods."
+ We focus below on Sigual-to-Noise (SNR) based detection at a single location as au example application., We focus below on Signal-to-Noise (SNR) based detection at a single location as an example application.
+ The SNR was the primary statistic used for detecting sources in high-cucrey astroplivsics before the introduction of maxiumuu-liselibood aud. wavelet-based methods., The SNR was the primary statistic used for detecting sources in high-energy astrophysics before the introduction of maximum-likelihood and wavelet-based methods.
+ Typically. E=3 owas used as the detection threshold. correspοiding to à=0.003 in the Gaussian regime.," Typically, $\frac{S}{N}=3$ was used as the detection threshold, corresponding to $\alpha=0.003$ in the Gaussian regime."
+ Hore we ayply the recipe in to derive an upper luit with SNR-based detectio1., Here we apply the recipe in \\ref{s:recipe} to derive an upper limit with SNR-based detection.
+ Our methods cau also be applie to more sophisticated «etection algorithius such as sliding-cell detection methods such as (11:uuden et 1198. Dobrzvcki ct 22000. Calderwood et 22001). and wavelet-based detection metlods such as (Darniani et 11997). (Vikhliuin ot 110997). (Freciuan et 2202). ete.," Our methods can also be applied to more sophisticated detection algorithms such as sliding-cell detection methods such as (Harnden et 1984, Dobrzycki et 2000, Calderwood et 2001), and wavelet-based detection methods such as (Damiani et 1997), (Vikhlinin et 1997), (Freeman et 2002), etc."
+ hupleimoeutatio1 of our technique for these methods will vary in detail. aud we leave these developnents for future work.," Implementation of our technique for these methods will vary in detail, and we leave these developments for future work."
+ We beein with a Caussian probability model for the source and background couuts (Step 1 in lpe]), We begin with a Gaussian probability model for the source and background counts (Step 1 in \\ref{s:recipe}) )
+electron-positron pair creation due to the electric Ποια of the electrosphere could be one of (he main observational signatures of quark stars.,electron-positron pair creation due to the electric field of the electrosphere could be one of the main observational signatures of quark stars.
+ It is the purpose of (he present paper to consider (he Schwinger process of pair creation in the electrosphere of the quark stars. by svstematically taking into account (he main physical characteristics of the environment in which pair production takes place.," It is the purpose of the present paper to consider the Schwinger process of pair creation in the electrosphere of the quark stars, by systematically taking into account the main physical characteristics of the environment in which pair production takes place."
+ The electric field in the electrosphere is not a constant. as assumed in the Schwinger model. but it is a rapidly decreasing function of the distance z from the quark star surface.," The electric field in the electrosphere is not a constant, as assumed in the Schwinger model, but it is a rapidly decreasing function of the distance $z$ from the quark star surface."
+ Therefore. in order to realistically describe the pair production process one must consider electron-positron creation in an inhomogeneous electric field.," Therefore, in order to realistically describe the pair production process one must consider electron-positron creation in an inhomogeneous electric field."
+ To study (he pair production process we adopt the tunnelling approach and the fermion production rate in the electric field of the electrosphere is derived., To study the pair production process we adopt the tunnelling approach and the fermion production rate in the electric field of the electrosphere is derived.
+ The production rate is a local quantity. it depends on the distance to the quark stars surface and il is a rapidly decreasing Function of z.," The production rate is a local quantity, it depends on the distance to the quark star's surface and it is a rapidly decreasing function of $z$."
+ The emissivity ancl enereyv [Iux due (o pair creation at the quark star surface is also considered., The emissivity and energy flux due to pair creation at the quark star surface is also considered.
+ An important factor which can reduce significantly the pair production rate is (the presence of a boundary (the quark star surface)., An important factor which can reduce significantly the pair production rate is the presence of a boundary (the quark star surface).
+ For electric Ποιά» perpendicular to the boundary. there is a significant reduction in (he magnitude of the pair production rate. which is also associated with an important qualitative change in the process. which becomes a periodic function of the distance to the quark star surface.," For electric fields perpendicular to the boundary there is a significant reduction in the magnitude of the pair production rate, which is also associated with an important qualitative change in the process, which becomes a periodic function of the distance to the quark star surface."
+ The present paper is organized as follows., The present paper is organized as follows.
+ In Section 2 we review (he basic properties of the electrosphere of quark stars., In Section 2 we review the basic properties of the electrosphere of quark stars.
+ The electron-positron rate production in (he electric field ol the electrosphere is obtained in Section 3., The electron-positron rate production in the electric field of the electrosphere is obtained in Section 3.
+ The boundary effects are considered in Section 4., The boundary effects are considered in Section 4.
+ In Section 5 we calculate the electron-positron pair flux of the electrosphere ancl compare our results with those obtained by Usov(1998a.b.2001).," In Section 5 we calculate the electron-positron pair flux of the electrosphere and compare our results with those obtained by \citet{Us98a,Us98b,Us01}."
+.. A brief summary of our results is given in Section 6., A brief summary of our results is given in Section 6.
+ In the electrosphere. electrons are held to the strange quark matter (5M) surface by an extremely strong electric field.," In the electrosphere, electrons are held to the strange quark matter (SQM) surface by an extremely strong electric field."
+ The thickness of the electrosphere is much smaller than the stellar radius Re109 em. and a plane-parallel approximation may be used to study its structure (Usovetal.2005).," The thickness of the electrosphere is much smaller than the stellar radius $R\simeq 10^6$ cm, and a plane-parallel approximation may be used to study its structure \citep{Us05}."
+. In Chis approximation all values depend only on the coordinate 2. Where (he axis z is perpendicular to (he SQAI surface (2= 0) and directed outward.," In this approximation all values depend only on the coordinate $z$, where the axis $z$ is perpendicular to the SQM surface $z=0$ ) and directed outward."
+ To lind the distributions of electrons and electric fields in (he vicinity of the SQAI surface. we use a simple Thomas-Fermi model considered by Aleocketal.(1986). and takeinto account (he finite temperature effects as discussed bv Nettierοἱal. (1995).," To find the distributions of electrons and electric fields in the vicinity of the SQM surface, we use a simple Thomas-Fermi model considered by \citet{Al86} and takeinto account the finite temperature effects as discussed by \citet{Ke95}. ."
+We construct a simple tov model of the galaxys star formation history (ΕΙ). iu the bulee CR<5 spc}. tuner disk (5«Rx15 kpc) and outer disk CR>15 kpc) of the galaxy.,"We construct a simple toy model of the galaxy's star formation history (SFH), in the bulge $R<5$ kpc), inner disk $5<R<15$ kpc) and outer disk $R>15$ kpc) of the galaxy."
+ Usine coustraiuts derived frou the galaxvs current star formation rate and stellar iuass surface density. we build a simple model that is consistent with observations in all three locations.," Using constraints derived from the galaxy's current star formation rate and stellar mass surface density, we build a simple model that is consistent with observations in all three locations."
+" Our observations show that the current star-forimation is uniformly distributed across the galaxy. πο we postulate that the overall SEIT cau be modeled by a period of coustaut star formation of some length that is the same evervwliere. superimposed on an older stellar population that dominates at the ceuter aud is virtually abseut at large τας,"," Our observations show that the current star-formation is uniformly distributed across the galaxy, so we postulate that the overall SFH can be modeled by a period of constant star formation of some length that is the same everywhere, superimposed on an older stellar population that dominates at the center and is virtually absent at large radii."
+ Our model incorporates the following constraints: A simple schematic of this model ΕΠ is shown iu Figue 10.., Our model incorporates the following constraints: A simple schematic of this model SFH is shown in Figure \ref{sfh}.
+" We then use Druzual Charlot (2003) models to predict D1000, iu the inner disk aud bulec. obtaining values of 1.3 and 1.5. respectively."," We then use Bruzual Charlot (2003) models to predict $D4000_n$ in the inner disk and bulge, obtaining values of 1.3 and 1.5, respectively."
+ These values are iu good agreement with the measured values of 1.1 and 1.6 in the corresponding radial bius., These values are in good agreement with the measured values of 1.4 and 1.6 in the corresponding radial bins.
+ Thus. through oulv a simple partition of mass into old aud voung conrponenuts. we can reasonably reproduce one of the kev spectral features of GASS35981.," Thus, through only a simple partition of mass into old and young components, we can reasonably reproduce one of the key spectral features of GASS35981."
+ We uote that varviug the formation time of the old component or doubliug its timescale to 2 Gar changes the results very little., We note that varying the formation time of the old component or doubling its timescale to 2 Gyr changes the results very little.
+ Likewise. adjusting the leneth of the more recent star-forming episode within the range 0.72 Car does not sienificautly affect our results.," Likewise, adjusting the length of the more recent star-forming episode within the range 0.7–2 Gyr does not significantly affect our results."
+ Finally. we can estimate the total fraction of the stellar nass added to the galaxy in the most recent star formation episode aud find that it is around20%.," Finally, we can estimate the total fraction of the stellar mass added to the galaxy in the most recent star formation episode and find that it is around."
+. Note. however. that at its current star formation rate. CGASS35981 will not exhaust its ITE reservoir for another NT Cr.," Note, however, that at its current star formation rate, GASS35981 will not exhaust its HI reservoir for another 5–7 Gyr."
+ We have concluded that a new episode of star formation began approximately lo Cyr ago in CASS35981L., We have concluded that a new episode of star formation began approximately 1 Gyr ago in GASS35981.
+ It is tempting to speculate that renewed star formation in this ealaxy is connected το the acquisition of its large III reservoir., It is tempting to speculate that renewed star formation in this galaxy is connected to the acquisition of its large HI reservoir.
+ So how did GCASS35981 acquire its gas?, So how did GASS35981 acquire its gas?
+ One possibility is that the gas was accreted as the result of au iuteraction with another eas-rich system., One possibility is that the gas was accreted as the result of an interaction with another gas-rich system.
+ CASS35981 does have two Iuimuinous ucielibors at projected distances of300 kpe aud 100 kpe with redshift differeuces of <d00l tie. CASS35981 does appear to be a 1ienber of a loose eroup.," GASS35981 does have two luminous neighbors at projected distances of 300 kpc and 400 kpc with redshift differences of $<400$ km $^{-1}$ —i.e., GASS35981 does appear to be a member of a loose group."
+ We do not. however. sce any sigus of an interaction or morecr in either the rotation curve or optical morphology of CASS35981.," We do not, however, see any signs of an interaction or merger in either the rotation curve or optical morphology of GASS35981."
+ Furthermore. a scenario in which GASS35981 is left with ος1010 AL. of HII after an encounter with either of its neighbors seenis fairly tuplausible.," Furthermore, a scenario in which GASS35981 is left with $2 \times 10^{10}$ $_\odot$ of HI after an encounter with either of its neighbors seems fairly implausible."
+ Simulations predict that no more than of the gas mass of auy donor galaxy should be stripped in an encounter (Dournaud. private communication). so woe would requie a donor galaxy with >1015 MIL. of HI!," Simulations predict that no more than of the gas mass of any donor galaxy should be stripped in an encounter (Bournaud, private communication), so we would require a donor galaxy with $>10^{11}$ $_\odot$ of HI!"
+ Such an encountercowed. however. be responsible for jostling an already-preseut III reservoir out of equilibrium. causing it to form stars.," Such an encounter, however, be responsible for jostling an already-present HI reservoir out of equilibrium, causing it to form stars."
+ If the eas was not acquired from another passing ealaxy. oue might speculate that CASS35981 accreted its III reservoir directlv from the siuvrouudius intergalactic medi at some point in the past.," If the gas was not acquired from another passing galaxy, one might speculate that GASS35981 accreted its HI reservoir directly from the surrounding intergalactic medium at some point in the past."
+ Since bliud III surveys such as ALFALFA (Ciovanclli et al., Since blind HI surveys such as ALFALFA (Giovanelli et al.
+ 2005) aud HIPASS (Barnes et al., 2005) and HIPASS (Barnes et al.
+ 2001) do not fiud dark UT clouds of 102 NE. the eas ist have eutered GASS35981 from au ionized phase if it entered all at once or over a short timescale.," 2001) do not find dark HI clouds of $10^{10}$ $_\odot$, the gas must have entered GASS35981 from an ionized phase if it entered all at once or over a short timescale."
+ Since its two identified ncighbors sueecst that CASS35981 resides iu a group. the presence of intri- gas could be fucling an unusually Ligh accretion rate onto CLASS35981.," Since its two identified neighbors suggest that GASS35981 resides in a group, the presence of intra-group gas could be fueling an unusually high accretion rate onto GASS35981."
+ Indeed. accretion from the," Indeed, accretion from the"
+so the solution to Eq.,so the solution to Eq.
+ 2) for a point mass AZ at 5—0 is the Newtonian potential ey=-GAL/r. and the consequent centripetal lorce obevs the inverse-square law.," \ref{peq} for a point mass $M$ at $r=0$ is the Newtonian potential $\psi_N=-GM/r$, and the consequent centripetal force obeys the inverse-square law."
+ addressed the disk velocity discrepancy for galaxies with his Mocified Newtonian Dynamics theory. MOND. for which Eq.," \citet{Mi:83a,Mi:83b} addressed the disk velocity discrepancy for galaxies with his Modified Newtonian Dynamics theory, MOND, for which Eq."
+" 1 is modified (BekensteinandMilgrom1954). with the replacement where the acceleration αμ is a universal constant and F’(r7)zzx [or e«1 and (a2)z1 [or r>>1. On setting e,210 “ems 7. the fanetional form of F obliges spiral galaxy disk velocily profiles to be generally flat. in agreement wilh observations."," \ref{eq:LN} is modified \citep{BeMi:84} with the replacement where the acceleration $a_0$ is a universal constant and ${\mathcal F}'(x^2)\approx x$ for $x \ll 1$ and ${\mathcal F}'(x^2)\approx 1$ for $x \gg 1.$ On setting $a_0\approx 10^{-8}$ cm $^{-2}$, the functional form of ${\mathcal F}$ obliges spiral galaxy disk velocity profiles to be generally flat, in agreement with observations."
+ BrownsteinandMoffat(2006) and MannheimaudO'Brien(2011) also devised non-Newtonian explanations [or the flat velocity profiles of spiral galaxies., \citet{BM:06} and \citet{MO:11} also devised non-Newtonian explanations for the flat velocity profiles of spiral galaxies.
+ Eckhardt.Pestana&Fischbach(2010) |EPF] considered non-Newtonian gravitation al (he scale of galaxy. superclusters., \citet{EPF:10} [EPF] considered non-Newtonian gravitation at the scale of galaxy superclusters.
+ For (he Lagrangian. the solution to 0£=0 is and (he solution to Eq.," For the Lagrangian, the solution to $\delta {\mathcal L}=0$ is and the solution to Eq."
+" 5. lor a point mass Af at r=0 has the form of a Yukawa potential (see Appendix). If gez» 0. this results in a Milne universe (hat is in full accord with cosmological expansion observations of (vpe Ia supernovae but. unlike the explication of Riessetal.(2007).. il requires neither dark matter (Q,,= 0) nor dark enerev (Q4= 0)."," \ref{yeq} for a point mass $M$ at $r=0$ has the form of a Yukawa potential (see Appendix), If $\mu>0$ , this results in a Milne universe that is in full accord with cosmological expansion observations of type Ia supernovae but, unlike the explication of \cite{Retal:07}, it requires neither dark matter $\Omega_m=0$ ) nor dark energy $\Omega_{\Lambda}=0$ )."
+" The choice pp|z5 Alpe then explains the scales of galaxy. superclusters and of the fundamental spectrum of the cosmic background radiation: ancl it explains why neighboring superclusters tend to be aligned in spongiform ""surfaces with vast empty regions between them.", The choice $\mu^{-1}\approx 5$ Mpc then explains the scales of galaxy superclusters and of the fundamental spectrum of the cosmic background radiation; and it explains why neighboring superclusters tend to be aligned in spongiform “surfaces” with vast empty regions between them.
+" ;The eravilon− mass corresponding− to jj1=5 \Mpe is− im=hy/e1.3.x10“3 >/ο-. so EPF conjectured that the mass of the (cosmological) graviton is m,eLO"" /c?. and that there could be heavier gravitons as well. but none that is lighter than m."," The graviton mass corresponding to $\mu^{-1}= 5$ Mpc is $m=\hbar \mu/c =1.3\times10^{-30}$ $^2$, so EPF conjectured that the mass of the (cosmological) graviton is $m_c\sim 10^{-30}$ $^2$, and that there could be heavier gravitons as well, but none that is lighter than $m_c$."
+ This led us to conjecture that a heavier (galaxy) mj graviton is responsible for the flat galaxy. velocity curves., This led us to conjecture that a heavier (galaxy) $m_g$ graviton is responsible for the flat galaxy velocity curves.
+" Anim, Yukawa potential cannot explain the fat curves. but Eckhardt(1993). had suggested that an i, exponential potential combined with a Newtonian potential could provide an explanation (see his Fie."," An $m_g$ Yukawa potential cannot explain the flat curves, but \cite{eck93} had suggested that an $m_g$ exponential potential combined with a Newtonian potential could provide an explanation (see his Fig."
+ 2)., 2).
+ Ilowever. histheoretical argument for (hie existence," However, histheoretical argument for the existence"
+with no further approximations and converting to non-LC coordinates vields what we may call à surrogate Hamiltonian for which the analvtie solution is exact.,with no further approximations and converting to non-LC coordinates yields what we may call a surrogate Hamiltonian for which the analytic solution is exact.
+ For timelike geodesies £ is small and for null geodesies 1/77 is a second order correction., For timelike geodesics $E$ is small and for null geodesics $1/r^2$ is a second order correction.
+ Thus to obtain a first order expression for error in the metric components we may neglect the term containing ££., Thus to obtain a first order expression for error in the metric components we may neglect the term containing $E$.
+ I it is bothersome that the Llamiltonian in equation (43)) is dependent on ££. one can eliminate it by a further mocification.," If it is bothersome that the Hamiltonian in equation \ref{hamilterr}) ) is dependent on $E$, one can eliminate it by a further modification."
+ Let us add a E/[Q[! term to the E Hamiltonian. which makes only a higher-order change to the solutions.," Let us add a $E/{|Q|^4}$ term to the $\Gamma$ Hamiltonian, which makes only a higher-order change to the solutions."
+" Introducing 42,4 by one finds which. when expanded in L/r to the appropriate order. is identical to (43)) less the term dependent on ££."," Introducing $H'_{\rm surr}$ by one finds which, when expanded in $1/r$ to the appropriate order, is identical to \ref{hamilterr}) ) less the term dependent on $E$ ."
+ The Hamiltonian (46)) is then the Hamiltonian for which equations (22)). (23)). and (24)) are exact solutions.," The Hamiltonian \ref{hamilterr2}) ) is then the Hamiltonian for which equations \ref{solnb}) ), \ref{solnu}) ), and \ref{solnull}) ) are exact solutions."
+ We have derived timelike and null geodesics in the leacing-orcder Schwarzschild metric in terms of elementary functions., We have derived timelike and null geodesics in the leading-order Schwarzschild metric in terms of elementary functions.
+ The expressions (22)) for bound orbits and (23)) for unbound orbits. together with (24)) for light rays. ave all simple eeneralizations of well-known expressions in. classical celestial mechanics.," The expressions \ref{solnb}) ) for bound orbits and \ref{solnu}) ) for unbound orbits, together with \ref{solnull}) ) for light rays, are all simple generalizations of well-known expressions in classical celestial mechanics."
+ Phe usual formulas for relativistic orbital precession and. light cellection are easily recovered., The usual formulas for relativistic orbital precession and light deflection are easily recovered.
+ A feature resembling the innermost stable circular orbit in the full Schwarzschild metric is also present., A feature resembling the innermost stable circular orbit in the full Schwarzschild metric is also present.
+ The technique we have used is a modification of the Levi-Civita or Ixustaanheimo-Stiefel regularization transformation and transforms the σουοσίο equation into a spherical harmonic oscillator., The technique we have used is a modification of the Levi-Civita or Kustaanheimo-Stiefel regularization transformation and transforms the geodesic equation into a spherical harmonic oscillator.
+ “Phe simplicity of the result. notwithstanding the non-trivial route used to derive it. hints at some underlying svmmetry in the Schwarzschild) problem.," The simplicity of the result, notwithstanding the non-trivial route used to derive it, hints at some underlying symmetry in the Schwarzschild problem."
+ We speculate that it is somehow related to the separability of the Llamilton-Jacobi and other equations in the Schwarzschild ancl Ixerr metrics but have not attempted to investigate this., We speculate that it is somehow related to the separability of the Hamilton-Jacobi and other equations in the Schwarzschild and Kerr metrics \citep[cf.][]{1983mtbh.book.....C} but have not attempted to investigate this.
+ As mentioned in the Introduction. the originalS motivation for this work was to find useful formulas applicablePI to the highlv-eccentrieIn Galactic-centre stars. whose orbits pass throughe a largeὃν rangeo of ogravitational regimes.," As mentioned in the Introduction, the original motivation for this work was to find useful formulas applicable to the highly-eccentric Galactic-centre stars, whose orbits pass through a large range of gravitational regimes."
+ὃν Future observations of these stars aiming to detect relativistic cllects will require computation of relativistic ellects on both stellar orbits and light ravs at many points along an orbit. for many orbits. in order to fit the orbital parameters.," Future observations of these stars aiming to detect relativistic effects will require computation of relativistic effects on both stellar orbits and light rays at many points along an orbit, for many orbits, in order to fit the orbital parameters."
+ The solutions in this paper allow a simpler. more cllicient method for carrying out those computations.," The solutions in this paper allow a simpler, more efficient method for carrying out those computations."
+ The analytic solutions will not be sullicient on their own because the CGalactic-centre stars also experience additional Newtonian perturbations due to local matter (2).. but they can be incorporated into numerical methods. specificallv. generalized leapfrog integrators.," The analytic solutions will not be sufficient on their own because the Galactic-centre stars also experience additional Newtonian perturbations due to local matter \citep{2008AJ....135.2398M}, but they can be incorporated into numerical methods, specifically, generalized leapfrog integrators."
+ Such algorithms evolve alternately uncer two Llamiltonians. which are integrable separately.," Such algorithms evolve alternately under two Hamiltonians, which are integrable separately."
+ Ehe idea goes back to 27. and ?.., The idea goes back to \cite{1991AJ....102.1528W} and \cite{1991CeMDA..50...59K}.
+ Some recent developments on adaptive stepsizes appear in ? and are applied to the specific problem of CGalactic-centre stars in 7.., Some recent developments on adaptive stepsizes appear in \cite{2007CeMDA..98..191E} and are applied to the specific problem of Galactic-centre stars in \cite{2009ApJ...703.1743P}.
+ We note. however. that the present work is limited to test particles. and hence will not be applicable for binary orbits or self-gravitating disc simulations unless a generalization is found.," We note, however, that the present work is limited to test particles, and hence will not be applicable for binary orbits or self-gravitating disc simulations unless a generalization is found."
+ Another potential application may be the use of the solutions in relativistic clise simulations as an alternative to the widelv used. pseudo-Newtonian potentials (seeespecially???).. an advantage being that the solutions in this paper are well-defined approximations andinclude a more complete repertoire of &eneral-relativistie elfects for the same computational budget.," Another potential application may be the use of the solutions in relativistic disc simulations as an alternative to the widely used pseudo-Newtonian potentials \citep[see
+especially][]{1980A&A....88...23P,1996ApJ...461..565A,2009A&A...500..213A}, an advantage being that the solutions in this paper are well-defined approximations andinclude a more complete repertoire of general-relativistic effects for the same computational budget."
+and were observed in the J aud UW bands with an Oll-uürslow suppressor spectrograph (OILS). mount ou the Subaru Telescope.,"and were observed in the J and H bands with an OH-airglow suppressor spectrograph (OH-S), mounted on the Subaru Telescope."
+ The spectral resolution (A/AA) is equal to 210 (J baud) and 120 (IT band)., The spectral resolution $\lambda/\Delta\lambda$ ) is equal to 210 (J band) and 420 (H band).
+" We also selected 2 of the 1 SDSS QSOs at 2~6 preseute> by(200L):: TD and I& observations were carried out with the Cooled Infrared Spectrograph iux Camera (CISCO) mounted on the Subaru telescope. with a spectral resolution of 210 and 330 respectively,"," We also selected 2 of the 4 SDSS QSOs at $z\sim6$ presented by: H and K observations were carried out with the Cooled Infrared Spectrograph and Camera (CISCO) mounted on the Subaru telescope, with a spectral resolution of 210 and 330 respectively."
+ We use the K-baud spectzuiiof the +=6.13 QSO J1118|5251 published by(2003)., We used the K-band spectrumof the $z=6.43$ QSO J1148+5251 published by.
+. Data were obtaiue« with the NIRSPEC spectrograph ou EKeckII. (spectra resolution of 1500)., Data were obtained with the NIRSPEC spectrograph on KeckII (spectral resolution of 1500).
+ We selected Lb of the 5 sources presented by at redshifts +>5.8., We selected 4 of the 5 sources presented by at redshifts $z>5.8$.
+ The K-band observations were carried out with VLT-ISAAC with a spectral resolution of 150., The K-band observations were carried out with VLT-ISAAC with a spectral resolution of 450.
+ We also included the Is-band spectra of JO303-0019 preseuted by(2009)., We also included the K-band spectrum of J0303-0019 presented by.
+. This faint :~6 QSO was selected in SDSS Stripe 82 and its spectrum was taken with VLT-ISAAC., This faint $z\sim6$ QSO was selected in SDSS Stripe 82 and its spectrum was taken with VLT-ISAAC.
+ Finally. we added 1 SDSS QSOs at 2>5.5 published by(2007).. who observed them with CNIRS on Comin South with a spectral resolution of 800.," Finally, we added 4 SDSS QSOs at $z>5.8$ published by, who observed them with GNIRS on Gemini South with a spectral resolution of 800."
+ We focus on the spectral region with restframe wavelengths 2000 iApese<3500À., We focus on the spectral region with restframe wavelengths $2000$ $< \lambda_{rest} < 3500$.
+. This region is characterized by the presence of the eemission lue. the underline non-stellay coutimmun. the Baluer pseudo contimmun aud the eeuissiou line forest.," This region is characterized by the presence of the emission line, the underlying non-stellar continuum, the Balmer pseudo continuum and the emission line forest."
+ The last three enussiou features. that are overlapped iu the spectral range of interest. are described iu the following sectionis.," The last three emission features, that are overlapped in the spectral range of interest, are described in the following sections."
+ The dominant component of a QSO spectrum is the uon-stellar coutimmiuu. modeled as a power-law: Typically. in our data. the determination of the slope coefficient a depends on the adopted fitting procedure aud on the observed spectral rauge.," The dominant component of a QSO spectrum is the non-stellar continuum, modeled as a power-law: Typically, in our data, the determination of the slope coefficient $\alpha$ depends on the adopted fitting procedure and on the observed spectral range."
+ Iu case of a wide wavelength coverage it is possible to choose fitting windows free of contributions bv other enüssion components., In case of a wide wavelength coverage it is possible to choose fitting windows free of contributions by other emission components.
+" For spectra with restricted wavelength coverage, the power-law contmnuuni has to be fitted sinultaueouslv with the other components. resulting iu a local estimate of the slope that mieht not be fully representative of the overall coutimmun shape of the QSO."," For spectra with restricted wavelength coverage, the power-law continuum has to be fitted simultaneously with the other components, resulting in a local estimate of the slope that might not be fully representative of the overall continuum shape of the QSO."
+ analyzed a sample of 96 QSOs at 2<3 by fitting the power-law continmiuu in 8 differeut windows free of strong features. and obtained a mean value for the slope of =1.3 with a 1-0 dispersion of 1.6.," analyzed a sample of 96 QSOs at $z<3$ by fitting the power-law continuum in 8 different windows free of strong features, and obtained a mean value for the slope of $-1.3$ with a $\sigma$ dispersion of 1.6."
+ estimated the local slope for a sauiple of ~ 100.000 SDSS QSOs at :«L95. fitting the the power-law plus au iron template to the wavelength rauge around four broad emission lines(IIo...ITJ.. aaud Iv).," estimated the local slope for a sample of $\sim$ 100.000 SDSS QSOs at $z<4.95$, fitting the the power-law plus an iron template to the wavelength range around four broad emission lines, and )."
+ In particular. from the analvsis of the waveleusth range adjacent the eenissiou line. they obtained a mean value for the local slope of —1.9 with a 1-6 dispersion of 0.8.," In particular, from the analysis of the wavelength range adjacent the emission line, they obtained a mean value for the local slope of $-1.3$ with a $\sigma$ dispersion of 0.8."
+ From our spectral decomposition we find cousisteut values (see Sec. 3.1))., From our spectral decomposition we find consistent values (see Sec. \ref{sec_fit}) ).
+ We model the Balmer pseudo-coutiuuun following(2003)., We model the Balmer pseudo-continuum following.
+".. We assmnue partially optically thick gas clouds with wniform temperature T,=15000 Ix. For wavelengths below the Baliner edge (Apr=3616 Aj). the Balmer spectrum can be parametrized as: where By(f..) is the Planck fiction at the electro- cluperature {νι τε is the optical depth at the BalimeLael edee =12007). aud. Foe Is the normalized flux density at the Baluer edge1982)."," We assume partially optically thick gas clouds with uniform temperature $T_e=15000$ K. For wavelengths below the Balmer edge $\lambda_{BE}=3646$ ), the Balmer spectrum can be parametrized as: where $B_{\lambda}(T_e)$ is the Planck function at the electron temperature $T_{e}$ , $\tau_{BE}$ is the optical depth at the Balmer edge =1, and $F_{norm}$ is the normalized flux density at the Balmer edge."
+.. The normalization. should ο deteriumned at Apese23675Α.. where no eeniüssion is present.," The normalization should be determined at $\lambda_{rest} \backsimeq 3675$, where no emission is present."
+ Since this waveleneth is either no— covered or has a very low S/N in our suuple. we fix the iorinalizatiou to a fraction of the coutimmun strength extrapolated at 3675 As F4=fp:Fhowertew(3675 A}.," Since this wavelength is either not covered or has a very low S/N in our sample, we fix the normalization to a fraction of the continuum strength extrapolated at 3675 : $F_{norm}=f_B\cdot F_{power-law}(3675$ $)$."
+ To define the relative streueth of the two compoucuts and monitor the effects ou the cestimate. we have run various tests with fp=O.1.0.3.0.5.0.8.1.," To define the relative strength of the two components and monitor the effects on the estimate, we have run various tests with $f_B=0.1, 0.3, 0.5, 0.8, 1$."
+ Since differcuces in the cestimates resulting from this test were less than measured errors (with fp ouly partially affecting the power-law normalization). we lave fixed fp=0.3 based ou the results by (2003).," Since differences in the estimates resulting from this test were less than measured errors (with $f_B$ only partially affecting the power-law normalization), we have fixed $f_B=0.3$ based on the results by ."
+. The uou emits a forest of lines. many of which are blended.," The ion emits a forest of lines, many of which are blended."
+ We fit the fforest using a modified version of the cuuissionLue template by, We fit the forest using a modified version of the emissionline template by
+only).,only).
+ Fig., Fig.
+ 5 shows that the red galaxies have similar colors to BC03 single-burst stellar population models at z=2.2 and formation redshifts of zform23., \ref{fig_cmd} shows that the red galaxies have similar colors to BC03 single-burst stellar population models at $z=2.2$ and formation redshifts of $z_{form}\ga3$.
+" Fitting BC03 models with exponentially declining star formation histories to the full SEDs confirms that the red galaxies are on average ~1 Gyr old (zform~3.3), contain little on-going star formation and span a stellar mass range of M,—0.1-5x10!!M."," Fitting BC03 models with exponentially declining star formation histories to the full SEDs confirms that the red galaxies are on average $\sim1$ Gyr old $z_{form}\sim3.3$ ), contain little on-going star formation and span a stellar mass range of $=0.1-5\times10^{11}M_{\odot}$ ."
+" In contrast,by blue overdensity galaxies —Πι«1.6) are best-fit ~0.1 Gyr BC03 models and(J; have masses M,—105—10!°Mo.(BCGs"," In contrast, blue overdensity galaxies $J_1-H_l<1.6$ ) are best-fit by $\sim0.1$ Gyr BC03 models and have masses $=10^8-10^{10}M_{\odot}$."
+") We find candidate *brightest cluster galaxies"" in each overdensity, with relatively large stellar masses: M—3x107, MP—1x10H, Mf=1x10!and"," We find candidate “brightest cluster galaxies” (BCGs) in each overdensity, with relatively large stellar masses: $M^A_{\ast}=3\times10^{11}$, $M^B_{\ast}=1\times10^{11}$, $M^C_{\ast}=1\times10^{11}$."
+" As shown in Fig. 4,,"," As shown in Fig. \ref{fig_sed},"
+" BCG A is a quiescent “red Ms. dead” galaxy, while BCGs B C may contain recent star formation."," BCG A is a quiescent “red and dead” galaxy, while BCGs B C may contain recent star formation."
+ A close inspection of BCGs A B in Fig., A close inspection of BCGs A B in Fig.
+" 2 suggests they also have distinct structural properties: the latter is (r,=270Skpc) while the former has a larger core =3*6 kpc) plus an extended diffuse stellar halo.", \ref{fig_composite} suggests they also have distinct structural properties: the latter is $r_e=2^{+0.3}_{-0.5}$ kpc)while the former has a larger core $r_e=3^{+0.5}_{-0.5}$ kpc) plus an extended diffuse stellar halo.
+" With(re 3—4 satellites within r~30 kpc, we speculate BCG B will undergo a series of mergers and perhaps *puff-up"" in size (?).."," With $3-4$ satellites within $r\sim30$ kpc, we speculate BCG B will undergo a series of mergers and perhaps “puff-up” in size \citep{hopkins_discriminating_2010}."
+" Clearly we are probing an epoch where even some of the most massive galaxies in the highest density regions were still forming a significant fraction of their stars (e.g., "," Clearly we are probing an epoch where even some of the most massive galaxies in the highest density regions were still forming a significant fraction of their stars \citep[e.g., ][]{glazebrook_high_2004,van_dokkum_star_2007,eisenhardt_clusters_2008,tran_reversal_2010}."
+"To help us interpret the z—2.2 overdensities, we will now characterize various global properties of the overdensities and compare them to known high-redshift galaxy clusters and protoclusters."," To help us interpret the $z=2.2$ overdensities, we will now characterize various global properties of the overdensities and compare them to known high-redshift galaxy clusters and protoclusters."
+" As shown in Fig. l1, "," As shown in Fig. \ref{fig_density}, ,"
+"the individual z=2.2 overdensities have spatial extents of r—30"" or 250 kpc.", the individual $z=2.2$ overdensities have spatial extents of $r=30\arcsec$ or $250$ kpc.
+ The galaxy clusters at z=1.6 show similar projected compact sizes (????)..," The galaxy clusters at $z\ga1.6$ show similar projected compact sizes \citep{andreon_jkcs_2009,papovich_spitzer-selected_2010,tanaka_spectroscopically_2010,gobat_mature_2011}."
+" Notably, the z=1.62 cluster (??) shows two galaxy clumps or subclusters over a region of ~1'."," Notably, the $z=1.62$ cluster \citep{papovich_spitzer-selected_2010,tanaka_spectroscopically_2010} shows two galaxy clumps or subclusters over a region of $\sim1$."
+. This is not unlike the configuration discussed here., This is not unlike the configuration discussed here.
+" In contrast, known protoclusters at z=2 are typically more diffuse, with lower overdensities and 8x larger sizes (??).."," In contrast, known protoclusters at $z\ga2$ are typically more diffuse, with lower overdensities and $\sim8\times$ larger sizes \citep{steidel_ly_2000,venemans_protoclusters_2007}. ."
+" Indeed, overdensities A B each show core surface densities of >50 galaxies arcmin~?, with 5—6 members at r<10""."," Indeed, overdensities A B each show core surface densities of $\ga50$ galaxies $^{-2}$, with $5-6$ members at $r\la10\arcsec$."
+" To estimate the total halo mass (Mhaio)) of each overdensity, we use the relation between aand aat z—2.2 from the halo occupancy distribution analysis of ? and apply these to the oof the central BCGs."," To estimate the total halo mass ) of each overdensity, we use the relation between and at $z=2.2$ from the halo occupancy distribution analysis of \citet{moster_constraints_2010} and apply these to the of the central BCGs."
+" Although the uncertainties involved in converting stellar mass to halo mass are significant, we that the overdensities have: MA,©6x1013, CrateMP,+ 1x10!,ME,=1x109M..."," Although the uncertainties involved in converting stellar mass to halo mass are significant, we estimate that the overdensities have: $M_{halo}^A\approx6\times10^{13}$, $M_{halo}^B\approx1\times10^{13}$, $M_{halo}^C\approx1\times10^{13}$."
+ The oof A is in the same rangeas estimates for the z=2.07 cluster , The of A is in the same rangeas estimates for the $z=2.07$ cluster \citep{gobat_mature_2011}. .
+Using a 1 Gpc? cosmological simulation (GiggleZ;(?)..Poole et al., Using a 1 $^3$ cosmological simulation (GiggleZ;Poole et al.
+",in prep.;",",in prep.;"
+" 2160? particles, WMAP5 cosmology, ?)), we find that z=2.2 simulated halos with these masses will grow into z=0 halos with mean masses of Μιωο~0.5—5x10!¢4Mo"," $2160^3$ particles, WMAP5 cosmology,\citealt{komatsu_five-year_2009}) ), we find that $z=2.2$ simulated halos with these masses will grow into $z=0$ halos with mean masses of $M_{halo}\sim0.5-5\times10^{14}$ ."
+these stars. and N (Herwig2003)..,"these stars, and $^{14}$ N \citep{herwig05}. ."
+ Extra N is produced in the 'C- N pocket. which is then converted to '°F in the following convective TP (seealsoGoriely&Mowlavi2000).," Extra $^{15}$ N is produced in the $^{13}$ $^{14}$ N pocket, which is then converted to $^{19}$ F in the following convective TP \citep[see
+also][]{goriely00}."
+. The introduction of a 'C-'N pocket of increases the 'F yield by605c.. while the C and N yields are unaffected.," The introduction of a $^{13}$ $^{14}$ N pocket of increases the $^{19}$ F yield by, while the C and N yields are unaffected."
+" In Model 3 of Table | we consider a nodel computed using the upper limit (UL) of the ""Fr. py! e rate."," In Model 3 of Table \ref{tab:uncertain_yields}
+ we consider a model computed using the upper limit (UL) of the $^{18}$ $\alpha,p$ $^{21}$ Ne rate."
+" This increases the ""F yield by a factor of ~2.5 (seedetailsanddiscus-sioninKarakasetal. 2008)."," This increases the $^{19}$ F yield by a factor of $\simeq$ 2.5 \citep[see details and
+discussion in][]{karakas08}."
+". When both a C-N pocket and the upper limit of the F. py! e reaction rate Is used (Model 4 of Table 1»). the '""F yield increases by a factor of three."," When both a $^{13}$ $^{14}$ N pocket and the upper limit of the $^{18}$ $\alpha,p$ $^{21}$ Ne reaction rate is used (Model 4 of Table \ref{tab:uncertain_yields}) ), the $^{19}$ F yield increases by a factor of three."
+" The other main nuclear uncertainties originate from the U Ctr.y) O and the ""Fo.ϱ) Νε reaction rates (seediscus-Lugaroetal. 2004)."," The other main nuclear uncertainties originate from the $^{14}$ $\alpha,\gamma$ $^{18}$ O and the $^{19}$ $\alpha,p$ $^{22}$ Ne reaction rates \citep[see discussion 
+in][]{lugaro04}."
+. For the latter. the latest evaluation by Ugaldeetal.(2008) needs to be tested in AGB models.," For the latter, the latest evaluation by \citet{ugalde08} needs to be tested in AGB models."
+ All models presented do not take account of the effects induced by carbon enhancement on the opacities of the cool external layers of AGB stars., All models presented do not take account of the effects induced by carbon enhancement on the opacities of the cool external layers of AGB stars.
+ When C/O > 1. C-bearing molecules. most notably C» and CN. increase the opacity of the external layers. causing the envelope to expand and the star to become larger and cooler (Marigo2002).," When C/O $>$ 1, C-bearing molecules, most notably $_2$ and CN, increase the opacity of the external layers, causing the envelope to expand and the star to become larger and cooler \citep{marigo02}."
+. Models of AGB stars of low mass and metallicity with C- and N-enhanced low-temperature opacities have been calculated using the Frascati Raphson Newton evolutionary code (FRANEC) code (Stranieroetal.2006;Cristallo2007).," Models of AGB stars of low mass and metallicity with C- and N-enhanced low-temperature opacities have been calculated using the Frascati Raphson Newton evolutionary code (FRANEC) code \citep{straniero06,cristallo07}."
+. In these models the mass-loss rate strongly increases with respect to models in which opacities are always calculated using the initial Z=107+ solar-scaled composition., In these models the mass-loss rate strongly increases with respect to models in which opacities are always calculated using the initial $Z = 10^{-4}$ solar-scaled composition.
+ The resulting yields (Model 5 of Table 1)) are = 5 times smaller than in the Karakas models., The resulting yields (Model 5 of Table \ref{tab:uncertain_yields}) ) are $\approx$ 5 times smaller than in the Karakas models.
+ For comparison. the results obtained with the same code. but using opacities calculated for the initial solar-scaled composition (Cristallo2006) are also reported in Table | (Model 6). and are in good agreement with the Karakas model. in spite of the different choices of mass-loss rate and treatment of the convective borders (seeStranieroetal.2006.fordetails)..," For comparison, the results obtained with the same code, but using opacities calculated for the initial solar-scaled composition \citep{cristalloThesis} are also reported in Table \ref{tab:uncertain_yields} (Model 6), and are in good agreement with the Karakas model, in spite of the different choices of mass-loss rate and treatment of the convective borders \citep[see][for details]{straniero06}."
+ It is evident that further work is required to address the uncertainties in the AGB fluorine yields at low metallicities., It is evident that further work is required to address the uncertainties in the AGB fluorine yields at low metallicities.
+ In particular. the inclusion of low-temperature carbon-enhanced opacities in the Karakas models (Karakas. Wood Campbell. in preparation). will provide an independent comparison to the results obtained by the FRANEC code and by Marigo (2002).," In particular, the inclusion of low-temperature carbon-enhanced opacities in the Karakas models (Karakas, Wood Campbell, in preparation) will provide an independent comparison to the results obtained by the FRANEC code and by \citet{marigo02}."
+. Since a clear dependence of the mass-loss rate on the metallicity has still not be identified. different. mass-loss prescriptions should be tested (Cristallo et al..," Since a clear dependence of the mass-loss rate on the metallicity has still not be identified, different mass-loss prescriptions should be tested (Cristallo et al.,"
+ in preparation)., in preparation).
+ Finally. we note that the possible occurrence of the dual shell flash at the beginning of the AGB phase may also affect fluorine production and needs to be investigated in detail (e.g.&Lattanzio 2008).," Finally, we note that the possible occurrence of the dual shell flash at the beginning of the AGB phase may also affect fluorine production and needs to be investigated in detail \citep[e.g,][]{campbell08}."
+. We have shown that it is possible to reproduce the C and F abundances observed in the CEMP-s star HE 130540132 via binary mass transfer from a companion by accretion of 11% of the mass lost by the primary star during its AGB phase., We have shown that it is possible to reproduce the C and F abundances observed in the CEMP-s star HE 1305+0132 via binary mass transfer from a companion by accretion of $\%$ of the mass lost by the primary star during its AGB phase.
+ The AGB star should have dredged-up at least 0.2 of its intershell material into the convective envelope by means of the TDU., The AGB star should have dredged-up at least $\simeq$ of its intershell material into the convective envelope by means of the TDU.
+ While rapidly rotating massive stars produce enough carbon and nitrogen to form CEMP stars. they do not appear to produce fluorine (Palaciosetal.2005).," While rapidly rotating massive stars produce enough carbon and nitrogen to form CEMP stars, they do not appear to produce fluorine \citep{palacios05}."
+. The binary formation scenario is thus favoured in the case of HE 130540132., The binary formation scenario is thus favoured in the case of HE 1305+0132.
+ In general. we predict that most CEMP stars formed by mass transfer from an AGB companion should also be FEMP stars. Le.. have [F/Fe] > «1.," In general, we predict that most CEMP stars formed by mass transfer from an AGB companion should also be FEMP stars, i.e., have [F/Fe] $>$ +1."
+ Hence. fluorine appears to be a useful discriminant between the different scenarios proposed for the formation of CEMP stars.," Hence, fluorine appears to be a useful discriminant between the different scenarios proposed for the formation of CEMP stars."
+ During the preparation of this manuscript. another halo object highly enriched in fluorine was discovered. the planetary nebula BoBn | (Otsukaetal.2008)..," During the preparation of this manuscript, another halo object highly enriched in fluorine was discovered, the planetary nebula BoBn 1 \citep{otsuka08b}."
+ The metallicity. as well as all the C and N abundances observed in this object are the same. within errors. as those of HE 1305+0132. which suggests that BoBn | has a close connection to CEMP stars. perhaps representing the evolutionary outcome of a CEMP star.," The metallicity, as well as all the C and N abundances observed in this object are the same, within errors, as those of HE 1305+0132, which suggests that BoBn 1 has a close connection to CEMP stars, perhaps representing the evolutionary outcome of a CEMP star."
+ On the other hand. the derived F abundance in this object is roughly a factor of three higher than that obtained for HE 1305+0132.," On the other hand, the derived F abundance in this object is roughly a factor of three higher than that obtained for HE 1305+0132."
+ This observation can be explained via the binary scenario only if we consider the F yields we computed including the *C- ΗΝ. pocket or the upper limit of the Fr.pj! Νο reaction.," This observation can be explained via the binary scenario only if we consider the F yields we computed including the $^{13}$ $^{14}$ N pocket or the upper limit of the $^{18}$$\alpha,p$ $^{21}$ Ne reaction."
+ Following this indication. we multiplied the F yields by a factor of three in our stellar population model and calculated a probability of for CEMP stars to have A('°F)> +4.75.," Following this indication, we multiplied the F yields by a factor of three in our stellar population model and calculated a probability of for CEMP stars to have $^{19}$ $ \geq +4.75$."
+ This result provides us with a possibility to alleviate the problem of the extremely small probability of the high F abundance assessed for HE 1305+0132., This result provides us with a possibility to alleviate the problem of the extremely small probability of the high F abundance assessed for HE 1305+0132.
+and verifving that the same period (within the uncertainty imits. described below) appears in every subset.,"and verifying that the same period (within the uncertainty limits, described below) appears in every subset."
+ We checked he significance of the period by a Bootstrap technique (Iron&Tibshirani1993) on a sample of 1000 artificial LCs constructed by randomly. permutating the magnitudes of each subset., We checked the significance of the period by a Bootstrap technique \cite{efr93} on a sample of 1000 artificial LCs constructed by randomly permutating the magnitudes of each subset.
+ Phe reality of the periodicity found in the D. V and Clear filters was found to be significant at a confidence level.," The reality of the periodicity found in the B, V and Clear filters was found to be significant at a confidence level."
+ The uncertainty in the value of the period was determined again by the Bootstrap technique. this time o» a least square fitting of a sine function to the data. subtracting it [rom the LC and then randomly. adding the residuals to the sine data values.," The uncertainty in the value of the period was determined again by the Bootstrap technique, this time by a least square fitting of a sine function to the data, subtracting it from the LC and then randomly adding the residuals to the sine data values."
+ Next we found the peak requeney of 1000 artificial LCs so created., Next we found the peak frequency of 1000 artificial LCs so created.
+ In the histogram of the distribution of the peak frequencies we determined he interval that contains 980 results., In the histogram of the distribution of the peak frequencies we determined the interval that contains 980 results.
+ This is given above as a significant interval around the period value., This is given above as a significant interval around the period value.
+ In order to check whether D416 varies on a time scale shorter than days. we analysed the data collected curing the three nights of continuous monitoring.," In order to check whether B416 varies on a time scale shorter than days, we analysed the data collected during the three nights of continuous monitoring."
+ All three nights cid not show any periodicitv nor any significant variability. i.e. above 15 mmag.," All three nights did not show any periodicity nor any significant variability, i.e. above 15 mmag."
+ When we compare the magnitude of D41I6 in Clear between 1987 and LOOT. we discover that in LOOT the star was fainter by 170.03 mag than in 1987.," When we compare the magnitude of B416 in Clear between 1987 and 1997, we discover that in 1997 the star was fainter by $\pm$ 0.03 mag than in 1987."
+ This comparison was done using three standard stars in the field of D416. which did not vary by more than 0.02 mag during the past Ll vears.," This comparison was done using three standard stars in the field of B416, which did not vary by more than 0.02 mag during the past 11 years."
+ We also measured the mean lux centred at the AGTIOO continuum of our APO spectra ancl used it to calibrate the stars V magnitude using the value of Fy at that wavelength given by Calzetti et al. (, We also measured the mean flux centred at the $\lambda$ 6100 continuum of our APO spectra and used it to calibrate the star's V magnitude using the value of $_{\lambda}$ at that wavelength given by Calzetti et al. (
+1995): we obtained a Vo magnitude of 16.7-0.1 mag lor D416. which is consistent with the value 16.76 (see 83.3) they give based on their observations of N33 uring the 1986 LOST seasons.,"1995); we obtained a V magnitude of $\pm$ 0.1 mag for B416, which is consistent with the value 16.76 (see 3.3) they give based on their observations of M33 during the 1986 1987 seasons."
+ Ht therefore seems that vw red. continuum. of D416 remained essentially constant xween the epochs of 1987 ancl 1998., It therefore seems that the red continuum of B416 remained essentially constant between the epochs of 1987 and 1998.
+ he facing in Clear is accordinglv due to a fading in the blue band of more wn 0.17 mae. since Clear covers the whole range of optical LIÉeht and our CCD is more sensitive in the red than in the Xue band.," The fading in Clear is accordingly due to a fading in the blue band of more than 0.17 mag, since Clear covers the whole range of optical light and our CCD is more sensitive in the red than in the blue band."
+ Another change in the D and V brightness of the star is apparent between the 1907-08 ancl 1998-99 seasons., Another change in the B and V brightness of the star is apparent between the 1997-98 and 1998-99 seasons.
+ In order to amplify the elfect we binned the cata into 7 jns in each season (roughly one month per bin) and then calculated the mean magnitude of each binned season., In order to amplify the effect we binned the data into 7 bins in each season (roughly one month per bin) and then calculated the mean magnitude of each binned season.
+ Fig., Fig.
+ displavs the binned D and. V LCS., \ref{BVLCbin} displays the binned B and V LCs.
+ One can see that in D the star brightened by 0.06 mag and in V it brightened w 0.04 mag between the two seasons., One can see that in B the star brightened by 0.06 mag and in V it brightened by 0.04 mag between the two seasons.
+ In the B case this xieghtening is significant. since the cdillerence in the mean magnitudes is larger than the combined standard deviation of cach season's mean. while in V. the uncertainty in the magnitude is a little larger than the dillerence in mean magnitudes.," In the B case this brightening is significant, since the difference in the mean magnitudes is larger than the combined standard deviation of each season's mean, while in V, the uncertainty in the magnitude is a little larger than the difference in mean magnitudes."
+ This again manilests the fact that much like the SÍ226 periodic oscillations. the longer term. variations are also occuring mainly in the blue end of the spectrum.," This again manifests the fact that much like the 26 periodic oscillations, the longer term variations are also occuring mainly in the blue end of the spectrum."
+ The calculated ephemeris for the D V filters is: Pons UID 2450672.2(+1.6)|8.26(+0.03) E The set of folded. LCs along with a fitted sinusoid are plotted in Fig., The calculated ephemeris for the B V filters is: $=$ HJD $2450672.2(\pm1.6)+8.26(\pm0.03)$ E The set of folded LCs along with a fitted sinusoid are plotted in Fig.
+ 5 and a plot of a binned B LC along with a two-harmonic fit is shown in Fig. 6.., \ref{LCfold} and a plot of a binned B LC along with a two-harmonic fit is shown in Fig. \ref{harmony}.
+ It is apparent in Fig., It is apparent in Fig.
+ 5. that the amplitude of the modulation. decreases," \ref{LCfold}
+ that the amplitude of the modulation decreases"
+curve with the spectroscopic period and epoch and used 0.03 phase bins to reveal if there is any optical variability due to orbital motion.,curve with the spectroscopic period and epoch and used 0.03 phase bins to reveal if there is any optical variability due to orbital motion.
+" Although the SNR of the binned light curve is low for detailed speculations, we compared the shape of the curve with the EW variations of hydrogen and helium lines We note that some studies suggest the mass of an O6.5V type star to be around 28 — 29 Me (Martinsetal.2005) rather than To infer the mass loss rate of the O-type star and the properties of the circumstellar matter, we determined the equivalent widths of each H and He line and also their variability during the orbit, which could be good indicators of physical processes taking place in the stellar wind."," Although the SNR of the binned light curve is low for detailed speculations, we compared the shape of the curve with the EW variations of hydrogen and helium lines We note that some studies suggest the mass of an O6.5V type star to be around 28 – 29 $M_{\odot}$ \citep{Martins2005} rather than To infer the mass loss rate of the O-type star and the properties of the circumstellar matter, we determined the equivalent widths of each H and He line and also their variability during the orbit, which could be good indicators of physical processes taking place in the stellar wind."
+" As mentioned in Section 3.2, we also measured the EWs of some other lines to check the value of interstellar reddening"," As mentioned in Section 3.2, we also measured the EWs of some other lines to check the value of interstellar reddening"
+stellar disk might explain the difference between the observed BFs and the simulated BFs (e.g.. Pierce&Slaughter. 1982.. Balthasar 1988 or Hadrava 2007)).,"stellar disk might explain the difference between the observed BFs and the simulated BFs (e.g., \citeauthor{pierce1982} \citeyear{pierce1982}, , \citeauthor{balthasar1988} \citeyear{balthasar1988} or \citeauthor{hadrava2006} \citeyear{hadrava2006}) )."
+ Using the broadening of about 1000 lines. we obtained macro-turbulence parameters (Ze7) of 3.4 km/s and 3.3 km/s for the primary and secondary. respectively.," Using the broadening of about 1000 lines, we obtained macro-turbulence parameters $\zeta_{RT}$ ) of 3.4 km/s and 3.3 km/s for the primary and secondary, respectively."
+" For comparison, the macro-turbulence parameter found for the Sun. às a disk- star. for weak and moderately strong lines is 3.0 km/s (Takeda.1995)."," For comparison, the macro-turbulence parameter found for the Sun, as a disk-integrated star, for weak and moderately strong lines is $\approx4.0$ km/s \citep{Takeda1995}."
+ The values obtained for the vsini of the secondary. using the two different methods described above. agree with each other and are also consistent with the literature value.," The values obtained for the $v sin i$ of the secondary, using the two different methods described above, agree with each other and are also consistent with the literature value."
+ Note that there might be a systematic difference between the obtained radial- of the secondary during the secondary eclipse and the fit to the radial-velocities (Fig. 7))., Note that there might be a systematic difference between the obtained radial-velocities of the secondary during the secondary eclipse and the fit to the radial-velocities (Fig. \ref{fig:beta_secondary}) ).
+ This might be caused by too high an uncertainty in the orbital parameters used during the data analysis in Section 3.1.., This might be caused by too high an uncertainty in the orbital parameters used during the data analysis in Section \ref{sect:center}.
+ The values obtained for the primary vs/ni do agree with the literature value within their «c error. but do not agree with each other within their [-σ errors.," The values obtained for the primary $v sin i$ do agree with the literature value within their $\sigma$ error, but do not agree with each other within their $\sigma$ errors."
+ The method applied in Section 3.1. uses the amplitude of the RM effect to derive the value for the projected rotational velocity., The method applied in Section \ref{sect:center} uses the amplitude of the RM effect to derive the value for the projected rotational velocity.
+ It assumes a linear limb-darkening coetficient., It assumes a linear limb-darkening coefficient.
+ As mentioned in the last paragraph. this might not be sufficient.," As mentioned in the last paragraph, this might not be sufficient."
+ Also in Section 3.2. a linear limb-darkening coefficient is used. but vs/ni is derived not only by theuse of the RM effect. but also by the shape of the BF outside of the eclipse.," Also in Section \ref{sect:shape} a linear limb-darkening coefficient is used, but $v sin i$ is derived not only by theuse of the RM effect, but also by the shape of the BF outside of the eclipse."
+ Therefore. we disregard the value for vsni obtained in Section 3.1 and take the value of 19.6 km/s from Section 3.2. as the projected rotational velocity for the primary star in CCyg.," Therefore, we disregard the value for $v sin
+i$ obtained in Section \ref{sect:center} and take the value of $19.6$ km/s from Section \ref{sect:shape} as the projected rotational velocity for the primary star in Cyg."
+ Due to the remaining mismatch between the data and the simulation. which might be due to the use of too simple a model. we consider the formal errors for Zgy and vsini as too small.," Due to the remaining mismatch between the data and the simulation, which might be due to the use of too simple a model, we consider the formal errors for $\zeta_{RT}$ and $v sin i$ as too small."
+ We would like to point out that the methods used are still in their infancy., We would like to point out that the methods used are still in their infancy.
+ Even with the derived values for the stellar radii. the orbital inclination. limb-darkening and differential rotation are not free of errors yet: these parameters are not normally accessible. or only with great difficulty. via spectroscopic data.," Even with the derived values for the stellar radii, the orbital inclination, limb-darkening and differential rotation are not free of errors yet; these parameters are not normally accessible, or only with great difficulty, via spectroscopic data."
+ The main focus of this work is the robust determination of the orientation of the stellar rotation axes., The main focus of this work is the robust determination of the orientation of the stellar rotation axes.
+ The values derived for the projection of the rotation axes onto the plane of the sky are given in the last two rows of Table 2.., The values derived for the projection of the rotation axes onto the plane of the sky are given in the last two rows of Table \ref{tab:fit}. .
+ The £s derived from the two different methods agree to within their errors despite the differentsystematic problems of each method., The $\beta$ s derived from the two different methods agree to within their errors despite the differentsystematic problems of each method.
+ Therefore. the," Therefore, the"
+Commonly found. in both open ancl globular clusters (GCs). blue stragelers (BSs) appear as an extension of the main-sequence (MS) in cluster colour-magnituce cliagranis (CMDs). occuvping the region that is just. brighter. and bluer than the main-sequence turn-olf. (MSTO) (Sandage 1953).,"Commonly found in both open and globular clusters (GCs), blue stragglers (BSs) appear as an extension of the main-sequence (MS) in cluster colour-magnitude diagrams (CMDs), occuyping the region that is just brighter and bluer than the main-sequence turn-off (MSTO) \citep{sandage53}."
+. BSs are thought to be produced. via the addition of hivelrogen to low-mass AIS stars (c.g.Sillsetal.2001:Lom-bardietal. 2002).," BSs are thought to be produced via the addition of hydrogen to low-mass MS stars \citep[e.g.][]{sills01, lombardi02}."
+. This can occur via multiple channels. most of which involve the mergers of low-mass MS. stars since a significant amount of mass is typically required. to reproduce the observed. locations of Bss in. CALDs (e.g.Sills&ανα 1999).," This can occur via multiple channels, most of which involve the mergers of low-mass MS stars since a significant amount of mass is typically required to reproduce the observed locations of BSs in CMDs \citep[e.g.][]{sills99}."
+. Stars in close binaries can merge if enough. orbital angular momentum is lost. which can be mediated by dvnamical interactions with other stars. magnetized stellar winds. tidal dissipation or even an outer triple companion2010).," Stars in close binaries can merge if enough orbital angular momentum is lost, which can be mediated by dynamical interactions with other stars, magnetized stellar winds, tidal dissipation or even an outer triple companion."
+. Alternatively. AIS stars can collide directly. although this is also thought to usually be mediated: by multiplestar systems (e.g.Leonard1989:&Livio1995:Fregeauctal.2004:Leigh&Sills 2011)..," Alternatively, MS stars can collide directly, although this is also thought to usually be mediated by multiplestar systems \citep[e.g.][]{leonard89, leonard95, fregeau04, leigh11b}."
+ First proposed by AMeCrea(1964)... BSs have also been hypothesized to form by mass-transfer from an evolving primary onto a normal AIS companion via Roche lobe overllow.," First proposed by \citet{mccrea64}, BSs have also been hypothesized to form by mass-transfer from an evolving primary onto a normal MS companion via Roche lobe overflow."
+ Despite numerous formation mechanisms having been proposed. a satisfactory explanation to account for the presence of BSs in star clusters eludes us still.," Despite numerous formation mechanisms having been proposed, a satisfactory explanation to account for the presence of BSs in star clusters eludes us still."
+ Whatever the dominant DS formation mechanism(s) operating in dense clusters. it is now thought to somehow involve multiple star systems.," Whatever the dominant BS formation mechanism(s) operating in dense clusters, it is now thought to somehow involve multiple star systems."
+ This was shown to be the case in even the dense cores of GC's (Leigh.Sills&Ixnigge2007.2008:Ixnigge. where collisions between single stars are thought to occur frequentlyA (Leonard1989).," This was shown to be the case in even the dense cores of GCs \citep{leigh07, leigh08,
+ knigge09} where collisions between single stars are thought to occur frequently \citep{leonard89}."
+. In Ixnigge.Leigh&Sills (2009).. we showed that the numbers of BSs in the cores of a large sample of GCs correlate with the core masses.," In \citet{knigge09}, we showed that the numbers of BSs in the cores of a large sample of GCs correlate with the core masses."
+ We argued that our results are consistent with what is expected if Bhs are descended from binary stars since this would. imply a dependence of the form Nes~fisM ous. where Nes is the number of BSs in the core. fj is the," We argued that our results are consistent with what is expected if BSs are descended from binary stars since this would imply a dependence of the form $N_{BS} \sim f_bM_{core}$ , where $N_{BS}$ is the number of BSs in the core, $f_b$ is the"
+enable us to build a clearer picture of the inner Galaxy as inferred from its elobular clusters.,enable us to build a clearer picture of the inner Galaxy as inferred from its globular clusters.
+ It is important to see whether globular clusters reseuble the field stellar population of the bulge and to what degree. if we wish to have a comprehensive formation picture of the bulee.," It is important to see whether globular clusters resemble the field stellar population of the bulge and to what degree, if we wish to have a comprehensive formation picture of the bulge."
+ Tere. we contine our previous work on absolute proper motions of globular clusters based on Soutlern Proper Motiou Program data.," Here, we continue our previous work on absolute proper motions of globular clusters based on Southern Proper Motion Program data."
+ We describe the nieasuremoenuts im Section 2. the derivation of the absolute proper motions im Section 3. and that of the velocities iu Section 1.," We describe the measurements in Section 2, the derivation of the absolute proper motions in Section 3, and that of the velocities in Section 4."
+ We discuss our results iu Section 5., We discuss our results in Section 5.
+ We continue our program to measure absolute proper motions of globular clusters as a part of the overall Southeru Proper Motion Program (SPM) described in a series of papers: Platais et al. (, We continue our program to measure absolute proper motions of globular clusters as a part of the overall Southern Proper Motion Program (SPM) described in a series of papers: Platais et al. (
+1998). Girard et al. (,"1998), Girard et al. ("
+1998). τάχα et al. (,"1998), Girard et al. ("
+2001. 2010).,"2004, 2010)."
+ The cluster part of the SPAL provided measurements for 25 clusters reported im a series of papers: Dinescu et al. (, The cluster part of the SPM provided measurements for 25 clusters reported in a series of papers: Dinescu et al. (
+1997. 1999ab.. 2003) aud Casetti-Dinescu et al. (,"1997, 1999ab, 2003) and Casetti-Dinescu et al. ("
+2007). nauued Papers through V. respectively.,"2007), named Papers I through V, respectively."
+ Together with the sample presentedI here. the SPM has measured 31 clusters. which represeuts 55% of the total number of clusters with suchmeasurements’.," Together with the sample presented here, the SPM has measured 34 clusters, which represents $55\%$ of the total number of clusters with such."
+ The first-epoch SPM survey was taken from 1965 to 1979. and is eutirelv photographic.," The first-epoch SPM survey was taken from 1965 to 1979, and is entirely photographic."
+ The SPM secoucd-epoch survey is approximately one third photographic. (taken from 1988 to 1998) aud two thirds CCD-based (taken from 2003 through 2008)., The SPM second-epoch survey is approximately one third photographic (taken from 1988 to 1998) and two thirds CCD-based (taken from 2003 through 2008).
+ The main survey CCD data started with a two-fold overlap of the frames (sce below)., The main survey CCD data started with a two-fold overlap of the frames (see below).
+ Later. single coverage was used. when it was found that a sinele CCD exposure was already astrometrically superior to the first-epoch plate material.," Later, single coverage was used, when it was found that a single CCD exposure was already astrometrically superior to the first-epoch plate material."
+ For the clusters. apart from the main survey. we take additional frames ceutercd ou cach cluster.," For the clusters, apart from the main survey, we take additional frames centered on each cluster."
+ In Table 1 we list the SPM field centers. aud the photographic aud CCD material used in the reductions. together with the epochs ofthe observatious.," In Table 1 we list the SPM field centers, and the photographic and CCD material used in the reductions, together with the epochs of the observations."
+ Iu three of our four fields preseuted here. the secoud-epoch positions are CCD-based.," In three of our four fields presented here, the second-epoch positions are CCD-based."
+ Iu the romiainiug fiele (SPAL 1136. cluster NGC GEL). the secoud-epoch positions are predominantly photographic. supplementcc by six CCD exposures taken ou the cluster itself.," In the remaining field (SPM 436, cluster NGC 6441), the second-epoch positions are predominantly photographic, supplemented by six CCD exposures taken on the cluster itself."
+ The three fields with no secoud-epoch photographic materia lacked. a sinall portion of CCD coverage from the main SPM survey., The three fields with no second-epoch photographic material lacked a small portion of CCD coverage from the main SPM survey.
+ These areas consist of a declination ba for field #578 amounting to 23% of the cutive area. two corners of field 4307 amounting to [€ of the entire area. and about 15 for #651 around the edges of the fiek (seo Table 1).," These areas consist of a declination band for field $\# 578$ amounting to $23\%$ of the entire area, two corners of field $\# 307$ amounting to $4\%$ of the entire area, and about $1\%$ for $\# 651$ around the edges of the field (see Table 1)."
+ For these limited areas. the secoud-cpoch positions are taken from ZTépparcos. Tvcho-2 (Perrvinau et al.," For these limited areas, the second-epoch positions are taken from $Hipparcos$, Tycho-2 (Perryman et al."
+ 1997). UCAC? (Zacharias et al.," 1997), UCAC2 (Zacharias et al."
+ 2000) aud 2\TASS point source catalog (Cutri et al., 2000) and 2MASS point source catalog (Cutri et al.
+ 2003)., 2003).
+ Given the construction of the iuput list of objects to be measured. these positions are predominantly frou 2MASS.," Given the construction of the input list of objects to be measured, these positions are predominantly from 2MASS."
+ The program clusters are at low Calactic latitude uid toward the bulee direction., The program clusters are at low Galactic latitude and toward the bulge direction.
+ Since extinction is liel oei these directions. we used Tipparcos stars to tic to an inertial refercuce svstem rather thin extragalactic objects.," Since extinction is high in these directions, we used $Hipparcos$ stars to tie to an inertial reference system rather than extragalactic objects."
+ Thus the proper iiotions are ou the Iuternational Celestial Reference System via. Lipparcos., Thus the proper motions are on the International Celestial Reference System via $Hipparcos$.
+ Propor-motion units throughout the paper are nias , Proper-motion units throughout the paper are mas $^{-1}$.
+The SPAL plates were taken with the 5l-cnu double astroeraph at Cesco Observatory in two passbands: blue (103:-O. no filter) and visual (LO8a-G | ΟΕ2159 filter).," The SPM plates were taken with the 51-cm double astrograph at Cesco Observatory in two passbands: blue (103a-O, no filter) and visual (103a-G + OG515 filter)."
+ The plate scale is 55.107/nuu. and cach field covers 6.57«6:37.," The plate scale is 55.1”/mm, and each field covers $6.3\arcdeg \times 6.3\arcdeg$."
+" The plates contain two exposures: one of 2 hours that reaches V.—I8 and an offset of 2 miuutes,", The plates contain two exposures: one of 2 hours that reaches $V\sim 18$ and an offset of 2 minutes.
+ Diving both exposures. an objective grating is used which produces a series of diffraction inuages ou either side of the ceutral. zero-order images.," During both exposures, an objective grating is used which produces a series of diffraction images on either side of the central, zero-order images."
+ The multiple sets of imagesμα for bright stars (for VLl. there are ouly long exposure. zero-order nuages) allow us to detect and model magnitude-cdependenut svstematics that affect positions aud cousequeutly proper motions.," The multiple sets of images for bright stars (for $V \ge 14$, there are only long exposure, zero-order images) allow us to detect and model magnitude-dependent systematics that affect positions and consequently proper motions."
+ The method has been described and thoroughly tested bv Cürard et al. (, The method has been described and thoroughly tested by Girard et al. (
+1998). and it shows that /JIipparcos stars with V—9 aud faint cluster stars (V.211) can be placed on a system that is largely free of svstematies.,"1998), and it shows that $Hipparcos$ stars with $V \sim 9$ and faint cluster stars $V \ge 14$ ) can be placed on a system that is largely free of systematics."
+ For very bright magnitudes (V.«8.9. the corrections are large and likewise their πουΤαλαΊος: thus it is possible that residual maenitude equation is left at the very bright end. aud for this reason. the very bright LTfpparcos stars are nof used.," For very bright magnitudes $V < 8-9$, the corrections are large and likewise their uncertainties; thus it is possible that residual magnitude equation is left at the very bright end, and for this reason, the very bright $Hipparcos$ stars are not used."
+ Since the magnitude correction is mace independently for each plate. and there are iltiple sunc-epochn plates (from 2 to 1) in each field. residual nagnitude equation errors will show up as scatter in the uultiple-plate data aud contribute to the formal proper-notion uncertainties derived frou positional scatter (sec Section 3).," Since the magnitude correction is made independently for each plate, and there are multiple same-epoch plates (from 2 to 4) in each field, residual magnitude equation errors will show up as scatter in the multiple-plate data and contribute to the formal proper-motion uncertainties derived from positional scatter (see Section 3)."
+ The plates are ieasured with the Yale PDS wicrodensitometer iu an object-by-object mode with a πο] size of 12.7 gan. For cach SPAL feld. we measure a preselected set of stars (see also Papers TW aud. V).," The plates are measured with the Yale PDS microdensitometer in an object-by-object mode with a pixel size of 12.7 $\mu$ m. For each SPM field, we measure a preselected set of stars (see also Papers IV and V)."
+ This set consists of all fPipparcos stars (typically 100 ver field a set of ~L800 Tyeho-2 stars (Perrvinan et al.," This set consists of all $Hipparcos$ stars (typically 100 per field), a set of $\sim 1800$ Tycho-2 stars (Perryman et al."
+ 1997). another). set of ~1000 bright (Πενεωc 13.5) UCAC?2 stars (Zacharias et al.," 1997), another set of $\sim 1000$ bright $R_{UCAC2} < 13.5$ ) UCAC2 stars (Zacharias et al."
+ 2000). reference stars and cluster-reeion stars.," 2000), reference stars and cluster-region stars."
+ Refercuce aud cluster stars are selected from UCAC? and 2MASS (Cutri et al., Reference and cluster stars are selected from UCAC2 and 2MASS (Cutri et al.
+ 2003) catalogs., 2003) catalogs.
+ Chister-reeion stars are selected to reside within a7’ radius from the cluster ceuter., Cluster-region stars are selected to reside within a $7\arcmin$ radius from the cluster center.
+ Refereuce stars consist of two subsets: one set uniformly distributed over the cutive field. the other as a vine around cach cluster witli radius 7/<rx11.," Reference stars consist of two subsets: one set uniformly distributed over the entire field, the other as a ring around each cluster with radius $7\arcmin \le r \le 11\arcmin$."
+ Within the ring aud cluster area; all objects found in UCAC? or 2\TASS are selected for the input list.," Within the ring and cluster area, all objects found in UCAC2 or 2MASS are selected for the input list."
+ For the uniformly distributed reference stars over the cutire field. we select a fraction of UCAC? ans 2ATASS stars imposed by the time lanitation of the PDS scans.," For the uniformly distributed reference stars over the entire field, we select a fraction of UCAC2 ans 2MASS stars imposed by the time limitation of the PDS scans."
+ For instance. for field #4578 where there are four clusters. we had a total iuput nuniber of ~3900 stars for the uniformly distributed reference stars. ~13600 for the ring stars. and ~11500 stars for cluster-reeiou stars.," For instance, for field $\# 578$ where there are four clusters, we had a total input number of $\sim 3900$ stars for the uniformly distributed reference stars, $\sim 13600$ for the ring stars, and $\sim 11500$ stars for cluster-region stars."
+ For field 4307. where there is one cluster. we bad a total input uuuber of ~12700 stars for the randomly distributed reference stars. ~7100 for the ring stars. aud ~3s800 for the cluster-xegiou stars.," For field $\#307$, where there is one cluster, we had a total input number of $\sim 12700$ stars for the randomly distributed reference stars, $\sim 7100$ for the ring stars, and $\sim 3800$ for the cluster-region stars."
+ The bright stars in the iuput catalog assure au appropriate maenitude range of various diffraction orders with which to model magnitude-cependeut systematics., The bright stars in the input catalog assure an appropriate magnitude range of various diffraction orders with which to model magnitude-dependent systematics.
+ Ouly LTipparcos stays are used to determine the, Only $Hipparcos$ stars are used to determine the
+"raclius ratio is bounced,",radius ratio is bounded.
+ We have shown that jets impacting an inclined ambient densitv gradient are subject to the development of internal waves and oblique shocks., We have shown that jets impacting an inclined ambient density gradient are subject to the development of internal waves and oblique shocks.
+ Ht is impractical to run the simulations toa point where the leading bow shock leaves (he computational domain. but we have shown that these internal structures are slow moving. and argued (hat their criving flows persist long alter the bow has passed the deflecting structure: thus an interesting question is the extent {ο which these internal structures would be observable in maps not dominated by the bow shock (which even if in proximity (to (he structures of interest. might. contribute little to the svnehrotvon emissivity — e.g.. Miocduszewskietal. (1997))).," It is impractical to run the simulations to a point where the leading bow shock leaves the computational domain, but we have shown that these internal structures are slow moving, and argued that their driving flows persist long after the bow has passed the deflecting structure; thus an interesting question is the extent to which these internal structures would be observable in maps not dominated by the bow shock (which even if in proximity to the structures of interest, might contribute little to the synchrotron emissivity – e.g., \citet{mhd}) )."
+" In the simulation wilh 5;=2.5. 6=65"" (see Figure 6)) the oblique. high pressure structure upstream of the flows head spans the jet in the plane parallel to the ambient density gradient. and in a plane orthogonal to that. constitutes a wedge of extent ~2.75 jet radii. and maximum width ~0.43 jet radii."," In the simulation with $\gamma=2.5$, $\phi= 65\arcdeg$ (see Figure \ref{defl25A}) ) the oblique, high pressure structure upstream of the flow's head spans the jet in the plane parallel to the ambient density gradient, and in a plane orthogonal to that, constitutes a wedge of extent $\sim 2.75$ jet radii, and maximum width $\sim 0.43$ jet radii."
+ It thus fills ~7% of the volume of the perturbed jet segment., It thus fills $\sim 7$ of the volume of the perturbed jet segment.
+" The peak pressure in (his structure is po;70.38. while in the adjacent cocoon per0.2. and in the initial state p;,;=0.05."," The peak pressure in this structure is $p_{peak}\sim 0.38$, while in the adjacent cocoon $p_{ext}\sim 0.2$, and in the initial state $p_{int}=0.05$."
+ If the emission is optically thin. with spectrum x»[°. there is equipartition between radiating particles and random magnetic field (ας~up). and the number of radiating particles is proportional to the internal enerev and (hus (he pressure. the emissivity is jxp.77? (Mioduszewskietal.1997)..," If the emission is optically thin, with spectrum $\propto \nu^{-\alpha}$ , there is equipartition between radiating particles and random magnetic field $u_e\sim
+u_B$ ), and the number of radiating particles is proportional to the internal energy and thus the pressure, the emissivity is $j\propto
+p^{\left(\alpha+3\right)/2}$ \citep{mhd}."
+ For a=0.5 the emissivity of the oblique structure is ου35 (mes that of the quiescent jet., For $\alpha=0.5$ the emissivity of the oblique structure is $\sim 35$ times that of the quiescent jet.
+ Allowing for the ~7% filling factor. (his structure would be 2.4 times brighter (han a corresponding segment of unperturbed jet.," Allowing for the $\sim 7$ filling factor, this structure would be $\sim 2.4$ times brighter than a corresponding segment of unperturbed jet."
+ The simulation with 5= and ó=35° (see Figure 11)) also shows a prominent oblique internal structure in the final time frame: again. the structure spans the jet in a plane parallel to the ambient gradient. and its section orthogonal to (his direction approximates an ellipse with major axis tilted slightly with respect to the jets inflow direction.," The simulation with $\gamma=5$ and $\phi= 35\arcdeg$ (see Figure \ref{defl50sA}) ) also shows a prominent oblique internal structure in the final time frame; again, the structure spans the jet in a plane parallel to the ambient gradient and its section orthogonal to this direction approximates an ellipse with major axis tilted slightly with respect to the jet's inflow direction."
+ The filling factor is ~504., The filling factor is $\sim 50$.
+.. The corresponding pressure values are Όρος©0.78. Perr00.2 and pi;=0.27 (the initial pressure being higher (han in the slow jet case lor fixed Mach number).," The corresponding pressure values are $p_{peak}\sim
+0.78$, $p_{ext}\sim 0.2$ and $p_{int}=0.27$ (the initial pressure being higher than in the slow jet case for fixed Mach number)."
+ Following the salle arguments as used above to estimate the component Πας. it follows that this structure would be ~3.2 times brighter than a corresponding segment of unperturbed jet.," Following the same arguments as used above to estimate the component flux, it follows that this structure would be $\sim 3.2$ times brighter than a corresponding segment of unperturbed jet."
+ We conclude that where jets are disturbed by ambient density. structure. jet emission may be enhanced by a sienilicant [actor (2— 3) locally. with little impact on the jets global dynamics.," We conclude that where jets are disturbed by ambient density structure, jet emission may be enhanced by a significant factor $2-3$ ) locally, with little impact on the jet's global dynamics."
+ A simple prescription for the appearance of the precessed jet (see Figure 12)) cannot be readily given. as the structure is highly asymmetric. andwhile the local jet boundary evolves al subrelativistic speed. significant Doppler boosting is associated wilh the instantaneous flow," A simple prescription for the appearance of the precessed jet (see Figure \ref{precA}) ) cannot be readily given, as the structure is highly asymmetric, andwhile the local jet boundary evolves at subrelativistic speed, significant Doppler boosting is associated with the instantaneous flow"
+the 6 worst outlying points (representing 1.25% of the full dataset) in order to obtain a fair measure of how well the template represents the bulk of the data.,the 6 worst outlying points (representing $1.25\%$ of the full dataset) in order to obtain a fair measure of how well the template represents the bulk of the data.
+ The reduced x? value reported in panel (c) is quite acceptable for the aligned data but clearly poor for the unaligned folded dataset., The reduced $\chi^2$ value reported in panel (c) is quite acceptable for the aligned data but clearly poor for the unaligned folded dataset.
+" Figure 9 shows a similar situation for a variable star with a period of 294 days located just 19 arcsec from the M31 core, corresponding to a projected distance from the core of just 72 pc for an M31 distance of 780 kpc."," Figure \ref{lc-comp2} shows a similar situation for a variable star with a period of 294 days located just 19 arcsec from the M31 core, corresponding to a projected distance from the core of just 72 pc for an M31 distance of 780 kpc."
+" In the unaligned case the peak in the periodogram is completely spurious, corresponding to the seasonal variation in average flux due to different parts of the variable phase being sampled in different observing seasons."," In the unaligned case the peak in the periodogram is completely spurious, corresponding to the seasonal variation in average flux due to different parts of the variable phase being sampled in different observing seasons."
+ Once again the reduced x? for the ΕΒΕ template is quite acceptable for the light- derived from aligned data but poor for the unaligned dataset., Once again the reduced $\chi^2$ for the RBF template is quite acceptable for the light-curve derived from aligned data but poor for the unaligned dataset.
+ The examples of Figures 8 and 9 show that careful photometric alignment is essential for the accurate characterisation of variable signals., The examples of Figures \ref{lc-comp} and \ref{lc-comp2} show that careful photometric alignment is essential for the accurate characterisation of variable signals.
+" It is noticeable that the first season of Angstrom data shows good consistency with subsequent seasons for the variable star shown in the right-hand column of Figure 8,, which is located 2/4 from the M31 core (about a half of the LT field width away)."," It is noticeable that the first season of Angstrom data shows good consistency with subsequent seasons for the variable star shown in the right-hand column of Figure \ref{lc-comp}, which is located $2 \farcm 4$ from the M31 core (about a half of the LT field width away)."
+ However the light-curve of the variable close to the M31 core shown in the right-hand column of Figure 9 is much more erratic during the first season than in subsequent seasons., However the light-curve of the variable close to the M31 core shown in the right-hand column of Figure \ref{lc-comp2} is much more erratic during the first season than in subsequent seasons.
+ Visual inspection of the images shows that the majority of images taken during the LT commissioning season in 2004 show evidence of scattered or reflected light close to the core., Visual inspection of the images shows that the majority of images taken during the LT commissioning season in 2004 show evidence of scattered or reflected light close to the core.
+ Improvements in light baffling made to the LT in early 2005 substantially improved the quality of later data., Improvements in light baffling made to the LT in early 2005 substantially improved the quality of later data.
+ Figure 10 shows six more examples of detected variable stars at angular distances from the M31 core ranging from 223 arcsec (near the edge of our reference field) down to just 9 arcsec., Figure \ref{examples} shows six more examples of detected variable stars at angular distances from the M31 core ranging from 223 arcsec (near the edge of our reference field) down to just 9 arcsec.
+" They illustrate good season-to-season consistency in the corrected photometry, as evidenced by the smooth continuity of the folded light-curve, and also reasonable characterisation of the noise, as indicated by the size of the reduced x? values with respect to the RBF template."," They illustrate good season-to-season consistency in the corrected photometry, as evidenced by the smooth continuity of the folded light-curve, and also reasonable characterisation of the noise, as indicated by the size of the reduced $\chi^2$ values with respect to the RBF template."
+ The variable in Figure 10 located at 9 arcsec from the core has a predictably very noisy light-curve and is at the threshold of detectability; whilst it displays clear periodic behaviour the structure of the variation is not well resolved due to noise., The variable in Figure \ref{examples} located at 9 arcsec from the core has a predictably very noisy light-curve and is at the threshold of detectability; whilst it displays clear periodic behaviour the structure of the variation is not well resolved due to noise.
+ We therefore determine that the technique works well at least beyond ~10 arcsec from the M31 core., We therefore determine that the technique works well at least beyond $\sim 10$ arcsec from the M31 core.
+" As Figure 6 shows, there are relatively few variable objects that are detectable above the noise so close to the M31 core but the examples in Figures 9 and 10 demonstrate that, where they are detectable, we are usually able to obtain reliable photometry even in this extreme high-background regime."," As Figure \ref{sigma-map} shows, there are relatively few variable objects that are detectable above the noise so close to the M31 core but the examples in Figures \ref{lc-comp2} and \ref{examples} demonstrate that, where they are detectable, we are usually able to obtain reliable photometry even in this extreme high-background regime."
+" In this paper we have shown that optimal difference imaging in regions of very high background levels, such as the bulge regions"," In this paper we have shown that optimal difference imaging in regions of very high background levels, such as the bulge regions"
+"vields QF.~[0.04(1.0/0.27)]20.15.and OFMeL—OFF0,85.","yields $\Omega^{\mathrm{FRW}}_{b} \sim [ 0.04 (1.0/0.27) ] 
+\simeq 0.15$,and $\Omega^{\mathrm{FRW}}_{\mathrm{DM}} 
+\simeq 1 - \Omega^{\mathrm{FRW}}_{b} \simeq 0.85$."
+ Estimating the various fractions of dark ancl barvonic matter which may currently be clumped versus unclumped (Bochner2011).. we choose the range V=(0.78.0.85.0.92.0.96.1.0) for our runs.," Estimating the various fractions of dark and baryonic matter which may currently be clumped versus unclumped \citep{BochnerAccelPaperI}, we choose the range $\Psi _{0} = (0.78,0.85,0.92,0.96,1.0)$ for our runs."
+ Four different zjyy values. five different V; values. and three different clumping evolution models gives us 4x5x3=60 simulation runs in (total.," Four different $z_\mathrm{init}$ values, five different $\Psi _{0}$ values, and three different clumping evolution models gives us $4 \times 5 \times 3 = 60$ simulation runs in total."
+" Residual Hubble diagrams of these model cosmologies show them to be broadly successful in reproducing aCosmological-Constant-like observed. acceleration. with the Va; runs resembling flat NCDM with 0.4. the Vii, runs resembling Q4~0.5—0.5. and (he Vj runs resembling 0.97."," Residual Hubble diagrams of these model cosmologies show them to be broadly successful in reproducing aCosmological-Constant-like observed acceleration, with the $\Psi _{\mathrm{Sqr}}$ runs resembling flat $\Lambda$ CDM with $\Omega_{\Lambda} \sim 0.3-0.4$ , the $\Psi _{\mathrm{Lin}}$ runs resembling $\Omega_{\Lambda} \sim 0.5-0.8$, and the $\Psi _{\mathrm{MD}}$ runs resembling $\Omega_{\Lambda} \sim 0.65-0.97$ ."
+ Qualitativelv. a ~cozen of these model cosmologies fit the data essentially as well as the best-fit flat ACDAM model.," Qualitatively, a $\sim$ dozen of these model cosmologies fit the data essentially as well as the best-fit flat $\Lambda$ CDM model."
+ Two panels containing several of these best runs are shown in Figures 2. and 3.., Two panels containing several of these best runs are shown in Figures \ref{FigLinPlot} and \ref{FigMDPlot}.
+ For our quantitative analysis below. we fit the residualdistance modulus function for each model. Apipes=Nn—Mp. to the SCP Union compilation of 307 Type la supernovae (SNe).," For our quantitative analysis below, we fit the residualdistance modulus function for each model, $\Delta \mu_{\mathrm{Pert}} \equiv \Delta (m - M)_{\mathrm{Pert}}$, to the SCP Union compilation \citep{KowalRubinSCPunion} of 307 Type Ia supernovae (SNe)."
+" Using their publicly available psx; ancl σαν data. we compute values of AZ, for these «Αμ Curves. minimized for eachrun by optimizing If""."," Using their publicly available $\mu _{\mathrm{SN}, i}$ and $\sigma _{\mathrm{SN}, i}$ data, we compute values of $\chi ^{2} _{\mathrm{Fit}}$ for these $\Delta \mu_{\mathrm{Pert}}$ curves, minimized for eachrun by optimizing $H^{\mathrm{Obs}}_{0}$."
+" Given these results. (he corresponding fit probability. values (£54) for eachrun are then computed using the \? distribution with Np,=(307—Nias) degrees of freedom. where New=3 for our V(/) models (optimizing V. zii. and LO). Now=2 for flat ACDM (optimizing Q4=1—Oy and 192""). and Nyq=1 for flat SCDM (optimizing £2!"" only)."," Given these results, the corresponding fit probability values $P_{\mathrm{Fit}}$ ) for eachrun are then computed using the $\chi ^{2}$ distribution with $N_{\mathrm{DoF}} = (307 - N_{\mathrm{Fit}})$ degrees of freedom, where $N_{\mathrm{Fit}} = 3$ for our $\Psi (t)$ models (optimizing $\Psi _{0}$, $z_\mathrm{init}$, and $H^{\mathrm{Obs}}_{0}$ ), $N_{\mathrm{Fit}} = 2$ for flat $\Lambda$ CDM (optimizing $\Omega _{\Lambda} \equiv 1 - \Omega _\mathrm{M}$ and $H^{\mathrm{Obs}}_{0}$ ), and $N_{\mathrm{Fit}} = 1$ for flat SCDM (optimizing $H^{\mathrm{Obs}}_{0}$ only)."
+" Devond just explaining (he apparent acceleration seen in the SNe data. a (rue model οἱ (he universe must satisfy a variety of complementary constraints before it can be regarded as a fullv-consistent replacement for the traditional ACDAM “Cosmic Concordance""."," Beyond just explaining the apparent acceleration seen in the SNe data, a true model of the universe must satisfy a variety of complementary constraints before it can be regarded as a fully-consistent replacement for the traditional $\Lambda$ CDM “Cosmic Concordance”."
+" To obtain a new concordance withoutDark Energy. we must relate (he observational (""dressed) parameters to the unperturbedFRW (“hare”) parameters of our moclels."," To obtain a new concordance withoutDark Energy, we must relate the observational (“dressed”) parameters to the unperturbedFRW (“bare”) parameters of our models."
+ These mathematical relationships are derived in Bochner(2011): here we will simply quote the output results for our best-filting models., These mathematical relationships are derived in \citet{BochnerAccelPaperI}; here we will simply quote the output results for our best-fitting models.
+ The parameters studied include: KY (versus measured Hnbble Constant 97): the physical age of the universe. (9: the unperturbed cosmic density. OFTW: the observed CMD acoustic scale. /V*: the apparent. (overall) cosmic equation of state. «Y's and the observed jerk (or jolt) parameter. Jf.," The parameters studied include: $H^\mathrm{FRW}_{0}$ (versus measured Hubble Constant $H^\mathrm{Obs}_{0}$ ); the physical age of the universe, $t^\mathrm{Obs}_{0}$ ; the unperturbed cosmic density, $\Omega^\mathrm{FRW}_\mathrm{M}$ ; the observed CMB acoustic scale,$l^\mathrm{Obs}_{\mathrm{A}}$ ; the apparent (overall) cosmic equation of state, $w^\mathrm{Obs}_{0}$ ; and the observed jerk (or jolt) parameter, $j^\mathrm{Obs}_{0}$ ."
+ Observables 22h. i7. and j£7 are computed [rom the Taylor series expansion coellicients of the Iuminosity distance [unction. dipo(z0 7). for το—0 (see.e.g..Riessetal.2004)..," Observables $H^\mathrm{Obs}_{0}$ , $w^\mathrm{Obs}_{0}$ , and $j^\mathrm{Obs}_{0}$ are computed from the Taylor series expansion coefficients of the luminosity distance function, $d_{\mathrm{L}, \mathrm{Pert}}(z^\mathrm{Obs})$ , for $z^\mathrm{Obs} \rightarrow 0$ \citep[see, e.g.,][]{RiessGoldSilver}.."
+"target stars, so they would also allow for a more detailed analysis of the abundances of stars in the Hercules stream (following,2010), which could shed more light on radial mixing in the Galactic disk.","target stars, so they would also allow for a more detailed analysis of the abundances of stars in the Hercules stream <cit.>[following, which could shed more light on radial mixing in the Galactic disk."
+ The last two panels ofFigure show that the predicted line-of-sight velocity distribution depends strongly on the dynamical parameters of the bar., The last two panels of show that the predicted line-of-sight velocity distribution depends strongly on the dynamical parameters of the bar.
+ Changing the pattern speed of the bar throughRoig shifts the predicted distribution and in particular the location of the Hercules feature., Changing the pattern speed of the bar through shifts the predicted distribution and in particular the location of the Hercules feature.
+" While this shift at a particular location is degenerate with the shift due to the uncertainty in the local circular velocity, results from various locations can be combined to break this degeneracy and the local circular velocity is also already quite well constrained2009a)."," While this shift at a particular location is degenerate with the shift due to the uncertainty in the local circular velocity, results from various locations can be combined to break this degeneracy and the local circular velocity is also already quite well constrained."
+ Changing the strength of the bar changes the relative heights of the two main peaks in the predicted velocity distribution., Changing the strength of the bar changes the relative heights of the two main peaks in the predicted velocity distribution.
+" However, the main peak of the velocity distribution might not be as simple as the one predicted here due to other dynamical effects."," However, the main peak of the velocity distribution might not be as simple as the one predicted here due to other dynamical effects."
+" Integrated measures, such as the number of stars in the main peak and the number of stars in the Hercules feature, could mitigate this somewhat."," Integrated measures, such as the number of stars in the main peak and the number of stars in the Hercules feature, could mitigate this somewhat."
+" While predicting actual constraints would involve detailed simulations of the expected data, it is clear that constraints on the dynamical properties of the bar can be derived from the experiment proposed in this paper."," While predicting actual constraints would involve detailed simulations of the expected data, it is clear that constraints on the dynamical properties of the bar can be derived from the experiment proposed in this paper."
+" The astrometricGaia mission will measure the parallaxes and proper motions of up to one billion stars, most of which will be disk stars2001)."," The astrometric mission will measure the parallaxes and proper motions of up to one billion stars, most of which will be disk stars."
+" With parallax and proper motion accuracies down to 10 uas and mission-averaged line-of-sight velocity uncertainties smaller than 10km s! for many stars to 17th magnitude,Gaia will be able to probe the kinematics of the disk out to several kpc in all directions2008)."," With parallax and proper motion accuracies down to 10$\mu$ as and mission-averaged line-of-sight velocity uncertainties smaller than 10km $^{-1}$ for many stars to 17th magnitude, will be able to probe the kinematics of the disk out to several kpc in all directions."
+" Crucially, it will provide large samples of giants in many of the promising regions inFigure2, with velocity measurements accurate enough to verify the predictions laid out in this paper."," Crucially, it will provide large samples of giants in many of the promising regions in, with velocity measurements accurate enough to verify the predictions laid out in this paper."
+" This will also allow the dynamical properties of the bar to be tightly constrained, something that will be hard forGaia to do using direct observations of the bulge due to the large extinction toward the Galactic center 2005)."," This will also allow the dynamical properties of the bar to be tightly constrained, something that will be hard for to do using direct observations of the bulge due to the large extinction toward the Galactic center ."
+"of 7 equally sized redshift bins, the median performance measures are plotted.","of 7 equally sized redshift bins, the median performance measures are plotted."
+ The magnitude-limited training set experiences improvements in performance for both the lowest and highest redshifts., The magnitude-limited training set experiences improvements in performance for both the lowest and highest redshifts.
+ The purity of the Type Ia sample found using the magnitude-limited training sample improves significantly at the lowest and highest redshifts by incorporating redshifts., The purity of the Type Ia sample found using the magnitude-limited training sample improves significantly at the lowest and highest redshifts by incorporating redshifts.
+" T'his highlights the tremendous value that host redshifts have in Type Ia classification, especially at breaking degeneracies between supernova type and redshift."," This highlights the tremendous value that host redshifts have in Type Ia classification, especially at breaking degeneracies between supernova type and redshift."
+distributions of points inside total secondary eclipse and (2) a subset of the part outside secondary eclipse. with the same number of points as (1) and with a randomly selected starting point.,"distributions of points inside total secondary eclipse and (2) a subset of the part outside secondary eclipse, with the same number of points as (1) and with a randomly selected starting point."
+" In the construction of the distributions. to minimize the effect of an arbitrarily selected histogram bu size. We used r""andom values of the sizes of the histogram bin. between Ix10°° and 7x10°° in units of normalized flux."," In the construction of the distributions, to minimize the effect of an arbitrarily selected histogram bin size, we used random values of the sizes of the histogram bin, between $\times$ $^{-6}$ and $\times$ $^{-6}$ in units of normalized flux."
+ We repeated this test using 500 different starting points and values of the bin. and we estimated the I-c errors as above.," We repeated this test using 500 different starting points and values of the bin, and we estimated the $\sigma$ errors as above."
+ This test resulted in a depth of 0.006640., This test resulted in a depth of $\pm$.
+0020%.. We adopted a final value of the secondary eclipse depth of 0.00620.002% from the results of the different tests., We adopted a final value of the secondary eclipse depth of $\pm$ from the results of the different tests.
+ The secondary eclipse appears slightly offset. at à |-σ level. from the expected center for a circular orbit. which could be interpreted às à possible non-zero eccentricity of ecos w=0.0140.01.," The secondary eclipse appears slightly offset, at a $\sigma$ level, from the expected center for a circular orbit, which could be interpreted as a possible non-zero eccentricity of $e\cos\omega$ $\pm$ 0.01."
+ This eccentricity is within the 1-c uncertainty of the orbital solution obtained from the radial velocity measurements (0.03+40.03 in Alonsoetal. 2008))., This eccentricity is within the $\sigma$ uncertainty of the orbital solution obtained from the radial velocity measurements $\pm$ 0.03 in \citealt{alo08a}) ).
+ In the optical. the observed flux of the planet can be expressed: Le.. the Foyretlectedsum of a Fooreeiittedcomponent of reflected light. a thermal re-emission of the incident stellar flux. and a thermal emission not related to the incident flux from the star (e.g.. internal heat or emission due to tidal forces).," In the optical, the observed flux of the planet can be expressed: i.e., the sum of a component of reflected light, a thermal re-emission of the incident stellar flux, and a thermal emission not related to the incident flux from the star (e.g., internal heat or emission due to tidal forces)."
+" The reflected light can be estimated as Fin=Ae(&y where A, is the geometric albedo."," The reflected light can be estimated as $F_{\rm p,reflected}=A_g\left(R_p\over a\right)^2$ where $A_g$ is the geometric albedo."
+ We can calculate the thermal component of the planet's emission (the combination of Fyyeonincd and Fina) using the same procedure as in Alonsoetal.(2009b).," We can calculate the thermal component of the planet's emission (the combination of $F_{\rm p,reemitted}$ and $F_{\rm p,internal}$ ) using the same procedure as in \cite{alo09b}."
+. For that purpose. we used a value of Tyyp=56254120 K for the star (Alonso2008) and the model spectrum of a GS8V star from Pickles(1998)..," For that purpose, we used a value of $_{\rm eff}$ $\pm$ 120 K for the star \citep{alo08a} and the model spectrum of a G8V star from \cite{pickles}."
+ We calibrated the model spectrum to obtain the same integrated flux as a Planck function. with the Το of the star., We calibrated the model spectrum to obtain the same integrated flux as a Planck function with the $_{\rm eff}$ of the star.
+" Taking the CoRoT spectral response function into account (Auvergneetal..2009).. the ratio of areas of the star and the planet. and the incident flux lost in the reflection (ApFy, where Ay is the Bond albedo). we searched for the brightness temperature that matched the observed eclipse depth."," Taking the CoRoT spectral response function into account \citep{auv09}, the ratio of areas of the star and the planet, and the incident flux lost in the reflection $A_{\rm B} F_\star$ where $A_{\rm B}$ is the Bond albedo), we searched for the brightness temperature that matched the observed eclipse depth."
+ We assumed a black-body emission for the planet and considered different values of the albedo., We assumed a black-body emission for the planet and considered different values of the albedo.
+ We assumed Ag=Ay. a reasonable assumption since more than of the stellar flux is emitted at the CoRoT wavelengths.," We assumed $A_{\rm B}=A_{\rm g}$, a reasonable assumption since more than of the stellar flux is emitted at the CoRoT wavelengths."
+" As an example. Ay~A, within for the four solar system giant planets when integrating A, between 0.4 and um (Karkoschka.1994)."," As an example, $A_{\rm B}\sim A_{\rm g}$ within for the four solar system giant planets when integrating $A_{\rm g}$ between 0.4 and $\,\mu$ m \citep{kar94}."
+.. The different possible solutions are plotted in Figure 6.., The different possible solutions are plotted in Figure \ref{fig:fig6}.
+" The zero-albedo brightness temperature in the CoRoT bandpass calculated this way and including the uncertainties on the Ty resulted in Tycosi 1910705, K. If. we further assume that Frimemal=O and that the planet is in thermalequilibrium!.. this temperature favors high values of the re-distribution factor f=0.60+40.14 that are greater than the maximum expected for no re-distribution from the day to the night sides (f= 0.5)."," The zero-albedo brightness temperature in the CoRoT bandpass calculated this way and including the uncertainties on the $_{\rm eff}$ resulted in $_{\rm p,CoRoT}$ $^{+90}_{-100}$ K. If we further assume that $F_{\rm p,internal}=0$ and that the planet is in thermal, this temperature favors high values of the re-distribution factor $f$ $\pm$ 0.14 that are greater than the maximum expected for no re-distribution from the day to the night sides $f=0.5$ )."
+ This may be explained by (1) departures from a blackbody planetary emission. (ii) a non-zero albedo. or (iii) emission of significant internal energy from the planet.," This may be explained by (i) departures from a blackbody planetary emission, (ii) a non-zero albedo, or (iii) emission of significant internal energy from the planet."
+ We defer the first possibility to future studies of radiative transfer in this planet., We defer the first possibility to future studies of radiative transfer in this planet.
+ Assuming blackbody radiation and negligible Fpintomat- Fig.," Assuming blackbody radiation and negligible $F_{\rm p,internal}$, Fig."
+ 6 shows that likely values of the albedo (corresponding to f between 0.25 and 0.5) are A.=0.06+0.06. which confirms the very low reflection of the planet obtained theoretically(e.g. Marleyetal.1999:Sudarsky2000:SeagerHoodetal.2008)).," \ref{fig:fig6} shows that likely values of the albedo (corresponding to $f$ between 0.25 and 0.5) are $A_{\rm g}=0.06\pm 0.06$ , which confirms the very low reflection of the planet obtained theoretically(e.g. \citealt{marley99,sud00,sea00,hood08}) ),"
+ and observationally: upper limits for the albedos of exoplanets in similar conditions as CoRoT-2b have been reported by Charbonneau et al. (, and observationally: upper limits for the albedos of exoplanets in similar conditions as CoRoT-2b have been reported by Charbonneau et al. (
+"1999. A,<0.3 for 7 Bootis b). Leight et al. (","1999, $A_g<$ 0.3 for $\tau$ Bootis b), Leight et al. ("
+"2003a.b. A,<0.39 for r Bootis b. A,«0.12 for HD 753289). by Roweetal.(2008) on HD 209458b (a 3«» upper limit of 0.17).","2003a,b, $A_g<$ 0.39 for $\tau$ Bootis b, $A_g<$ 0.12 for HD 75289), by \cite{rowe} on HD 209458b (a $\sigma$ upper limit of 0.17)."
+" Finally. CoRoT-Ib was found to be such that A,«0.20 using CoRoT’s red channel of the light curve (Snellenetal..2009).. and independently by Alonsoetal.(2009b) using the white channel."," Finally, CoRoT-1b was found to be such that $A_{\rm g}<0.20$ using CoRoT's red channel of the light curve \citep{sne09}, and independently by \cite{alo09b} using the white channel."
+ To distinguish. among the different components of the observed planetary flux. we would need chromatic information.," To distinguish among the different components of the observed planetary flux, we would need chromatic information."
+ The CoRoT data are delivered in three band-passes. but we could not reach the level of precision needed to detect the secondary in these data.," The CoRoT data are delivered in three band-passes, but we could not reach the level of precision needed to detect the secondary in these data."
+ We attribute this to the difficulty of achieving a good correction of the jitter in the satellites pointing., We attribute this to the difficulty of achieving a good correction of the jitter in the satellite's pointing.
+ As a concluding remark. the depth of the secondary eclipse of CoRoT-2b ts roughly the same as a transit of a planet in frontof a solar radius star.," As a concluding remark, the depth of the secondary eclipse of CoRoT-2b is roughly the same as a transit of a planet in frontof a solar radius star."
+ The significance of the detection of the secondary eclipse thus emphasizes the excellent capabilities, The significance of the detection of the secondary eclipse thus emphasizes the excellent capabilities
+"In both events discussed here the Doppler-flow indicates symmetrical flows of ~40 km/s. It is unclear whether the 100 km/s components of EE1 are Doppler flows, since they can also be interpreted as blends by the lines at 120.613 nm and 120.704 nm (as discussed above).","In both events discussed here the Doppler-flow indicates symmetrical flows of $\approx$ 40 km/s. It is unclear whether the 100 km/s components of EE1 are Doppler flows, since they can also be interpreted as blends by the lines at 120.613 nm and 120.704 nm (as discussed above)."
+" In case of an interpretation as Doppler flow, this would imply a multi-component event with two sources in the slit area."," In case of an interpretation as Doppler flow, this would imply a multi-component event with two sources in the slit area."
+" Because of the ambiguity however, we do not discuss this issue any further."," Because of the ambiguity however, we do not discuss this issue any further."
+ The Doppler-flow pattern seen in the line profiles is almost unchanged in magnitude of the line shift and in the location along the slit., The Doppler-flow pattern seen in the line profiles is almost unchanged in magnitude of the line shift and in the location along the slit.
+ This excludes any lateral movement exceeding ~500 km along or across the slit within the 3-minute duration of the event., This excludes any lateral movement exceeding $\approx$ 500 km along or across the slit within the 3-minute duration of the event.
+" Within 3 minutes however, a bi-directional jet moving at 40 km/s should have reached 7,200 km in each component."," Within 3 minutes however, a bi-directional jet moving at 40 km/s should have reached 7,200 km in each component."
+ This requires that the direction of the jet deviates less than 1? from the LOS., This requires that the direction of the jet deviates less than $1\degree$ from the LOS.
+ Such a scenario is very unlikely., Such a scenario is very unlikely.
+ It is simply not possible that in such an event upflow and downflow stay over minutes stationary within the wide slit., It is simply not possible that in such an event upflow and downflow stay over minutes stationary within the wide slit.
+" Also the fact, that no increase in size along the slit is observed is a strong argument against a linear moving, collimated jet."," Also the fact, that no increase in size along the slit is observed is a strong argument against a linear moving, collimated jet."
+" A similar argument holds for EE2, that lasts even longer."," A similar argument holds for EE2, that lasts even longer."
+" The conflict of lacking apparent motion is so evident in the examples shown here, that we now adopt the alternative flow configuration."," The conflict of lacking apparent motion is so evident in the examples shown here, that we now adopt the alternative flow configuration."
+" If we assume a spicule-like feature that is as narrow as as the spectrometer slit and crosses the slit at some angle below 90, then the redshifted portion and the blueshifted portion will appear simultaneously in spectroscopic data and can stay without apparent motion for an extended period of time, exactly as observed in EE1 and in EE2."," If we assume a spicule-like feature that is as narrow as as the spectrometer slit and crosses the slit at some angle below $90\degree$, then the redshifted portion and the blueshifted portion will appear simultaneously in spectroscopic data and can stay without apparent motion for an extended period of time, exactly as observed in EE1 and in EE2."
+" The double-peak in the emission of EE] and EE2 may also be an effect of the spinning motion, if we assume the repetition of a brightness maximum after completing a full revolution after «200 s. Alternatively, the double-peaked lightcurve may have something to do with the occurrence of double-threaded jets that have been reported from XRT observations (Kamioetal., 2010)."," The double-peak in the emission of EE1 and EE2 may also be an effect of the spinning motion, if we assume the repetition of a brightness maximum after completing a full revolution after $\approx$ 200 s. Alternatively, the double-peaked lightcurve may have something to do with the occurrence of double-threaded jets that have been reported from XRT observations \citep{Kamio10}."
+. Motivated by the plausible solution of the old discrepancy we looked for suitable candidates of solar phenomena as conceivable counterparts for our double-component EEs., Motivated by the plausible solution of the old discrepancy we looked for suitable candidates of solar phenomena as conceivable counterparts for our double-component EEs.
+" Such candidates could be type II spicules or Rapid Blueshifted Events (RBEs) (DePontieuetal.,2009,2011;McIntoshal,2009;RouppevanderVoort,2009) since they have very similar characteristics in terms of velocity, lifetime, size, and repeatability."," Such candidates could be type II spicules or Rapid Blueshifted Events (RBEs) \citep{DePontieu09,DePontieu11,Scott09,vdVoort09} since they have very similar characteristics in terms of velocity, lifetime, size, and repeatability."
+ A direct proof of helicity in RBEs by imaging instruments has to our knowledge not been reported yet and may be difficult to achieve., A direct proof of helicity in RBEs by imaging instruments has to our knowledge not been reported yet and may be difficult to achieve.
+" There is, however, indirect evidence, since rotation in macrospicules — believed to be bundles of substructures — was often observed."," There is, however, indirect evidence, since rotation in macrospicules – believed to be bundles of substructures – was often observed."
+" On the small side, evidence for helicity in regular spicules has been reported in literature as already mentioned."," On the small side, evidence for helicity in regular spicules has been reported in literature as already mentioned."
+" Fig.3 shows several spicules and a macrospicule in a SUMER raster in obtained on August 18, 1996 in a coronal hole location."," Fig.3 shows several spicules and a macrospicule in a SUMER raster in obtained on August 18, 1996 in a coronal hole location."
+ The panels show a brightness raster (left) and a dopplergram (right)., The panels show a brightness raster (left) and a dopplergram (right).
+" It is obvious that the macrospicule is spinning like a bended cylinder, but there is no signature of this motion in the spectroheliogram."," It is obvious that the macrospicule is spinning like a bended cylinder, but there is no signature of this motion in the spectroheliogram."
+" This demonstrates that even in such large structures imagers are principally unable to observe the rotation, unless finestructures can be resolved."," This demonstrates that even in such large structures imagers are principally unable to observe the rotation, unless finestructures can be resolved."
+" Similar observations of ""tornados' have been reported by Pike&Ma-son(1998) from SoHO-CDS data.", Similar observations of 'tornados' have been reported by \citet{Pike98} from -CDS data.
+ We use the observed helicity in a macrospicule — a much larger feature than the EEs discussed here — together with the published premise of no obvious distinction between macrospicules and other spicules as support for our argument., We use the observed helicity in a macrospicule – a much larger feature than the EEs discussed here – together with the published premise of no obvious distinction between macrospicules and other spicules as support for our argument.
+ The cartoon in Fig., The cartoon in Fig.
+ 4 shows typical SUMER line profiles calculated for a spiraling spicule at various aspect angles., 4 shows typical SUMER line profiles calculated for a spiraling spicule at various aspect angles.
+" We assume three components of the total radiance, two from the spinning motion with a tangential velocity of +60 km/s that contribute with each."," We assume three components of the total radiance, two from the spinning motion with a tangential velocity of $\pm$ 60 km/s that contribute with each."
+ As the third component we assume a faint flow of 100 km/s along the spicule — typical for type-II spicules — that contributes with5%., As the third component we assume a faint flow of 100 km/s along the spicule – typical for type-II spicules – that contributes with.
+". The angle between spicule (or upflow) and LOS, 6, is set as 0°, 30°, 45°, 90°, 135°, 150°, and 180? for the seven cases (the cases with 0>90? are mirror symmetric to cases (A) to (C) and not shown in Fig."," The angle between spicule (or upflow) and LOS, $\theta$, is set as $\degree$ , $\degree$, $\degree$, $\degree$, $\degree$, $\degree$, and $\degree$ for the seven cases (the cases with $\theta > 90\degree$ are mirror symmetric to cases (A) to (C) and not shown in Fig."
+ 4)., 4).
+ The spectrum in case (C) is very similar to those presented here and can quantitatively reproduce our observations., The spectrum in case (C) is very similar to those presented here and can quantitatively reproduce our observations.
+ In many observations the red and the blue component of the EE are observed together with a component at rest., In many observations the red and the blue component of the EE are observed together with a component at rest.
+ We attribute this zero-velocity component to the background emission of the solar disk that is also visible in optical thin emission., We attribute this zero-velocity component to the background emission of the solar disk that is also visible in optical thin emission.
+" In our case, however, the foreground emission of the EEs is so much stronger that it outshines the background."," In our case, however, the foreground emission of the EEs is so much stronger that it outshines the background."
+" The fact that in SUMER spectra no EEs are observed in coronal lines is — besides the fact that no really good coronal lines for disk observations exist in the SUMER wavelength range — not in contradiction to the scenario suggested by DePontieuetal. (2011), who found that the heated volume is outside the leading edge of the jet."," The fact that in SUMER spectra no EEs are observed in coronal lines is – besides the fact that no really good coronal lines for disk observations exist in the SUMER wavelength range – not in contradiction to the scenario suggested by \citet{DePontieu11}, who found that the heated volume is outside the leading edge of the jet."
+" If heating takes place whilethe jet propagates and expands, then spectrometers, in"," If heating takes place whilethe jet propagates and expands, then spectrometers, in"
+between equivalent stars with the same helium abundance.,between equivalent stars with the same helium abundance.
+ This has the effect of making the star brighter. even if the material is not mixed efficiently (2).," This has the effect of making the star brighter, even if the material is not mixed efficiently ."
+. Therefore. collisions involving stars from different populations should be taken into consideration when constructing the blue straggler colour and luminosity functions in clusters with multiple populations.," Therefore, collisions involving stars from different populations should be taken into consideration when constructing the blue straggler colour and luminosity functions in clusters with multiple populations."
+" Here we explore the results of such a calculation,", Here we explore the results of such a calculation.
+ We use the STARS code ($22) to calculate the evolution tracks of both the progenitor stars and the collision products., We use the STARS code \ref{sec:stars}) ) to calculate the evolution tracks of both the progenitor stars and the collision products.
+ The structure of the collision products themselves is calculated using the MMAS code ($27 ))., The structure of the collision products themselves is calculated using the MMAS code \ref{sec:mmas}) ).
+" The parameter space for our models consists of the masses of he two colliding stars. Av; and M». the time of collision /,,j and he helium abundance of the colliding stars. Υ and Y»."," The parameter space for our models consists of the masses of the two colliding stars, $M_1$ and $M_2$, the time of collision $t_\mathrm{coll}$ and the helium abundance of the colliding stars, $Y_1$ and $Y_2$."
+ In principle he impact parameter for the collision spans another dimension in parameter space. but we will limit ourselves to non-rotating collision models here.," In principle the impact parameter for the collision spans another dimension in parameter space, but we will limit ourselves to non-rotating collision models here."
+ In essence. we assume that the collision oroduet has an efficient way to lose excess angular momentum by the action of a magnetic field) and that rotational mixing is unimportant.," In essence, we assume that the collision product has an efficient way to lose excess angular momentum by the action of a magnetic field) and that rotational mixing is unimportant."
+ With this assumption. the effect of a non-zero impact »urameter on the structure of the collision product is small and can be ignored.," With this assumption, the effect of a non-zero impact parameter on the structure of the collision product is small and can be ignored."
+ We calculated a grid of models that maps out the relevant yortion of parameter space for blue straggler formation in globular clusters., We calculated a grid of models that maps out the relevant portion of parameter space for blue straggler formation in globular clusters.
+" Our computational grid covers a mass range of 0.4...0.8M, for both stars."," Our computational grid covers a mass range of $0.4\ldots
+0.8\mathrm{M}_\odot$ for both stars."
+ We will refer to the most massive star as the primary and the least massive one as the secondary., We will refer to the most massive star as the primary and the least massive one as the secondary.
+ To account for the evolution of the parent stars as well as the evolution of the collision product since time of collision the collision time is varied between 6000Myr and 12000Myr in steps of |000Myr.," To account for the evolution of the parent stars as well as the evolution of the collision product since time of collision the collision time is varied between $6\,000 \,\mathrm{Myr}$ and $12\,000 \,\mathrm{Myr}$ in steps of $1\,000 \,\mathrm{Myr}$."
+ The heavy element content Z is set to Z=0.001 for most of our models and abundances are sealed to solar values., The heavy element content $Z$ is set to $Z = 0.001$ for most of our models and abundances are scaled to solar values.
+" The initial (ZAMS) helium abundance in the stars is set to Y=Y,42Z. where Y, is 0.24.0.32 or 0.40."," The initial (ZAMS) helium abundance in the stars is set to $Y = Y_0 + 2 Z$, where $Y_0$ is $0.24, 0.32$ or $0.40$."
+ This parameter range was chosen to cover both normal helium abundances CY;= 0.24) and the most extreme helium abundance proposed in the literature (Y=0.40. in NGC 2808. 22)).," This parameter range was chosen to cover both normal helium abundances $Y_0 = 0.24$ ) and the most extreme helium abundance proposed in the literature $Y_0 = 0.40$, in NGC 2808, )."
+ The initial hydrogen abundance X=|—YZ and the primary and secondary are allowed to have a different Yo., The initial hydrogen abundance $X = 1 - Y - Z$ and the primary and secondary are allowed to have a different $Y_0$.
+ The range of parameters in our grid is summarised in Table [.., The range of parameters in our grid is summarised in Table \ref{tab:grid}.
+ The model sets A. B and C all involve just collisions between stars of the same helium content.," The model sets A, B and C all involve just collisions between stars of the same helium content."
+" Model set D also includes collisions between stars with Y=0.24 and Y,=0.32 and is a superset of models A and B. Similarly model set E is à superset of model sets A. B and C that also contains the cross-collisions involving stars with Y;=0.40."," Model set D also includes collisions between stars with $Y_0 = 0.24$ and $Y_0 = 0.32$ and is a superset of models A and B. Similarly model set E is a superset of model sets A, B and C that also contains the cross-collisions involving stars with $Y_0 = 0.40$."
+ Model sets A. B and D have been calculated for Z=0.0003 as well as Z=0.001.," Model sets A, B and D have been calculated for $Z = 0.0003$ as well as $Z = 0.001$."
+ The structure of our collision products was calculated using the Make Me A Star (MMAS) code by?.., The structure of our collision products was calculated using the Make Me A Star (MMAS) code by.
+ MMAS approximates the structure of a collision product using an algorithm known as entropy sorting., MMAS approximates the structure of a collision product using an algorithm known as entropy sorting.
+ The idea behind this algorithm is that the quantity A=p/p? which is related to the thermodynamic entropy. increases (nearly) monotonically inside stars from centre to surface.," The idea behind this algorithm is that the quantity $A = p/\rho^{5/3}$, which is related to the thermodynamic entropy, increases (nearly) monotonically inside stars from centre to surface."
+ In the absence of strong shocks. as is the case for low-velocity collisions. A is conserved in the fluid elements from both colliding stars.," In the absence of strong shocks, as is the case for low-velocity collisions, $A$ is conserved in the fluid elements from both colliding stars."
+ The structure of the collision product can then be approximated by sorting the mass shells of both parent stars in order of increasing A., The structure of the collision product can then be approximated by sorting the mass shells of both parent stars in order of increasing $A$.
+ As a result of stellar evolution A decreases in the core. so that the core of the collision product is most likely to resemble the core of the most evolved parent star(22).," As a result of stellar evolution $A$ decreases in the core, so that the core of the collision product is most likely to resemble the core of the most evolved parent star."
+ Our evolutionary models are calculated using a version of Eggleton's stellar evolution code(22)... hereafter STARS.," Our evolutionary models are calculated using a version of Eggleton's stellar evolution code, hereafter STARS."
+" The STARS code solves the equations of stellar structure and. thenuclear energy generation rate simultaneously on an adaptive non- non-Eulerian (""Eggletonian"") grid(2).", The STARS code solves the equations of stellar structure and thenuclear energy generation rate simultaneously on an adaptive non-Lagrangian non-Eulerian (“Eggletonian”) grid.
+". Since we have update the nuclear reaction rates to the recommended values from(21. with the exception of the '""Ntp.30O reaction. for which we use the recommended rate from and?."," Since we have update the nuclear reaction rates to the recommended values from, with the exception of the $^{14}\mathrm{N} (\mathrm{p}, \gamma)^{15}\mathrm{O}$ reaction, for which we use the recommended rate from and."
+. We use the opacity tables of2.. which combine the OPAL opacities from?.. the low temperature molecular opacities from?.. electron scattering opacities from and the conductive opacities from?.," We use the opacity tables of, which combine the OPAL opacities from, the low temperature molecular opacities from, electron scattering opacities from and the conductive opacities from."
+. The assumed heavy-element composition is scaled to the solar mixture of(2)., The assumed heavy-element composition is scaled to the solar mixture of.
+. Chemical mixing due to convection and thermohaline mixing is taken into account., Chemical mixing due to convection and thermohaline mixing is taken into account.
+ All models are computed with a mixing-length ratio //Hp=2.0., All models are computed with a mixing-length ratio $l/H_P=2.0$.
+ As in. we have neglected convective overshooting in all models deseribed here.," As in, we have neglected convective overshooting in all models described here."
+ Its effect on the evolution of stars in the mass range of our collision products (0.8M.. — 1.6M.} is negligible(2)., Its effect on the evolution of stars in the mass range of our collision products $0.8 \mathrm{M}_\odot$ – $1.6 \mathrm{M}_\odot$ ) is negligible.
+. On the giant branch and the early AGB we adopt the Reimers- mass loss prescription from??.. scmasslos," On the giant branch and the early AGB we adopt the Reimers-like mass loss prescription from, M = (1 + )."
+"sM = Because we canno calculate through the core helium flash. we stop the code at the onset of the flash and construct a ""zero-age horizontal branch"" model with the correct total mass. core mass and composition."," Because we cannot calculate through the core helium flash, we stop the code at the onset of the flash and construct a “zero-age horizontal branch” model with the correct total mass, core mass and composition."
+ This construction is performed by pseudo-evolving a low mass core burning helium star., This construction is performed by pseudo-evolving a low mass core burning helium star.
+ Composition changes due to helium burning are disabled during this process. but hydrogen is allowed to burn normally so that the helium core can grow in mass.," Composition changes due to helium burning are disabled during this process, but hydrogen is allowed to burn normally so that the helium core can grow in mass."
+ Material of the appropriate envelope composition is accreted on the star until its mass matches that of the last pre-flash model., Material of the appropriate envelope composition is accreted on the star until its mass matches that of the last pre-flash model.
+ The helium core is then allowed to grow to the desired mass due to hydrogen. shell burning until the helium core mass too matches the mass in our pre-flash model., The helium core is then allowed to grow to the desired mass due to hydrogen shell burning until the helium core mass too matches the mass in our pre-flash model.
+ This procedure gives us a core helium burning star of the correct total mass. core mass and composition. but we ignore any mass loss or mixing that may arise as a result of the helium flash itself.," This procedure gives us a core helium burning star of the correct total mass, core mass and composition, but we ignore any mass loss or mixing that may arise as a result of the helium flash itself."
+ This is not an unreasonable assumption although some authors have found mixing during calculations of the helium flash at very low metallicity (22).., This is not an unreasonable assumption although some authors have found mixing during calculations of the helium flash at very low metallicity .
+ Recent 3D hydrodynamical calculations of the helium flash by are, Recent 3D hydrodynamical calculations of the helium flash by are
+cautioned. however. that these orbits are highly. superficial and it seems highly unlikely that such extreme orbits exist at the expense of more pedestrian ones in every case. although it may be true for one.,"cautioned, however, that these orbits are highly superficial and it seems highly unlikely that such extreme orbits exist at the expense of more pedestrian ones in every case, although it may be true for one."
+ Outside the tidal radius (at 1.9 Ape) circular orbits can survive indefinitely. since there exist no stars to scatter and drag the GC.," Outside the tidal radius (at $1.9~kpc$ ) circular orbits can survive indefinitely, since there exist no stars to scatter and drag the GC."
+ The only issue is the probability that the 4 large GC's can simultaneously have no projected radii greater han 1.4Ape. which we estimate to be greater than a third. which for a unique svstem is credible.," The only issue is the probability that the 4 large GCs can simultaneously have no projected radii greater than $1.4~kpc$, which we estimate to be greater than a third, which for a unique system is credible."
+ A final approach is o ignore the largest GC€ which greatly relieves the timing »oblem. since the other four CC's can withstand the DE for substantially longer interior to the tidal radius.," A final approach is to ignore the largest GC which greatly relieves the timing problem, since the other four GCs can withstand the DF for substantially longer interior to the tidal radius."
+ ‘Taking into account the discussion. above and the uncertainties of the analytical approach to DE in MOND. it is probably not worth attaching too much significance to the existence of these 5 GC's.," Taking into account the discussion above and the uncertainties of the analytical approach to DF in MOND, it is probably not worth attaching too much significance to the existence of these 5 GCs."
+ Phe bigeest mystery. for the DM paradigm is why this was seen as a critical problem in the first. place., The biggest mystery for the DM paradigm is why this was seen as a critical problem in the first place.
+ CAVA's research is supported. by the University of Torino and Regione Piemonte., GWA's research is supported by the University of Torino and Regione Piemonte.
+ Partial support from the INFN erant, Partial support from the INFN grant
+are the clleets of LIB stars on our integrated. Balmer indices in § 5.4.1..,are the effects of HB stars on our integrated Balmer indices in $\S$ \ref{subsubsec:HorizontalBranchStars}.
+ Returning to Figures S. to 12.. we find the agreement between the (οὐ. Mg 6 and Ales indices is good.," Returning to Figures \ref{fig:swbIVA} to \ref{fig:swbVII}, we find the agreement between the $\langle$ $\rangle$, Mg $b$ and $_2$ indices is good."
+ In general. the position of the clusters on the SSP grids using these different metallicity indicators are consistent within the uncertainties.," In general, the position of the clusters on the SSP grids using these different metallicity indicators are consistent within the uncertainties."
+ The LIS. Hop and Hep indices also behave similarly. although there are some indications of systematic dillerences between their age predictions for the vounger clusters.," The $\beta$, $\delta_{\rm F}$ and $\gamma_{\rm F}$ indices also behave similarly, although there are some indications of systematic differences between their age predictions for the younger clusters."
+ The SWB types lie in relatively tight groups on the SSP model erids., The SWB types lie in relatively tight groups on the SSP model grids.
+ As one goes to older S\WB types in Figures 8 to 12.. these groups trace a characteristic shape. moving from the top-left (voung ages. Fe/ll] 2 1.0). to centre-right (intermediateages. bef) 2 1.0). to bottoneleft (old ages. ΓΟΗ « 1.0).," As one goes to older SWB types in Figures \ref{fig:swbIVA} to \ref{fig:swbVII}, , these groups trace a characteristic shape, moving from the top-left (young ages, [Fe/H] $\gsim$ –1.0), to centre-right (intermediate ages, [Fe/H] $\gsim$ –1.0), to bottom-left (old ages, [Fe/H] $<$ –1.0)."
+ One cluster in Figure 12. clearly stands out as being significantly vounger than the other SWB VIL clusters. ancl provides a nice illustration of the advantage of integrated spectroscopy over photometry in separating age anc metallicity cllects.," One cluster in Figure \ref{fig:swbVII} clearly stands out as being significantly younger than the other SWB VII clusters, and provides a nice illustration of the advantage of integrated spectroscopy over photometry in separating age and metallicity effects."
+ Εμίν cluster. NGC 1865. has an SSP-derived. age of < 1.0 Gyr. rather than the ~ 10 Cir implied by its SWB tvpe. (," This cluster, NGC 1865, has an SSP-derived age of $<$ 1.0 Gyr, rather than the $\sim$ 10 Gyr implied by its SWB type. ("
+1997) obtained an age of 0.9 Cyr for this cluster from the position of its CALD turnoll. which is consistent with our value. (,"1997) obtained an age of 0.9 Gyr for this cluster from the position of its CMD turnoff, which is consistent with our value. ("
+1997) attributed its SWB nmis-classification to a combination of the stochastic elfects of bright stars in the cluster. and its relative faintness contrasted with a dense stellar background.,"1997) attributed its SWB mis-classification to a combination of the stochastic effects of bright stars in the cluster, and its relative faintness contrasted with a dense stellar background."
+ As our principle metallicity indicators we have chosen our best-measured magnesium dominant. Lick index (Meo). and mean of two iron indices (Fe5270 and be5335) in the form of Fe}.," As our principle metallicity indicators we have chosen our best-measured magnesium dominant Lick index $_2$ ), and mean of two iron indices (Fe5270 and Fe5335) in the form of $\langle$ $\rangle$."
+ As age-sensitive indices. we use 47 and KEEOO models). in addition to our. best-measured higher-order Balmer lines [+p and Hj models only).," As age-sensitive indices, we use $\beta$ and KFF99 models), in addition to our best-measured higher-order Balmer lines $\gamma_{\rm F}$ and $\delta_{\rm F}$ models only)."
+ Ages and metallicitics are obtained for each cluster by interpolating the model grids using a programmnie., Ages and metallicities are obtained for each cluster by interpolating the model grids using a programme.
+ Where clusters lie olf the grids. linear extrapolation is used.," Where clusters lie off the grids, linear extrapolation is used."
+ Uncertainties are derived by perturbing the line-strength indices by their corresponding measurement error., Uncertainties are derived by perturbing the line-strength indices by their corresponding measurement error.
+ Because of the non-orthogonal nature of the model erids. two cilfering uncertainties in age and two dillering uncertainties in metallicity are obtained.," Because of the non-orthogonal nature of the model grids, two differing uncertainties in age and two differing uncertainties in metallicity are obtained."
+ As discussed. previously. due to the effects of the LIB on the Balmer indices. the metal-poor. old regions models often have two age solutions for the SWB VIL clusters.," As discussed previously, due to the effects of the HB on the Balmer indices, the metal-poor, old regions models often have two age solutions for the SWB VII clusters."
+ In these cases. we adopt the prior that. these clusters are Galactic GC analogues. and adopt the older ages if the clusters are consistent with these old isochrones (see 8 5.4 for futher explanation.)," In these cases, we adopt the prior that these clusters are Galactic GC analogues, and adopt the older ages if the clusters are consistent with these old isochrones (see $\S$ \ref{subsec:AgeComparisons} for futher explanation.)"
+ Whilst we wish to compare the SSP model predictions ο derived ages and metallicitics. an important eature of the SSP models should be emphasised.," Whilst we wish to compare the SSP model predictions to derived ages and metallicities, an important feature of the SSP models should be emphasised."
+ Due to 1 non-orthogonality of the SSP grids. changes in the =wetallicity estimates of the clusters effect the derived. ages {the clusters and here is still an age-metallicity degeneracy.," Due to the non-orthogonality of the SSP grids, changes in the metallicity estimates of the clusters effect the derived ages of the clusters and – there is still an age-metallicity degeneracy."
+ As a direct result. of this. the IXEE99 models. which have a tencency ο systematically over-precict metallicities with respect to 10 models (see & 5.1)). svstematically under-predict the ‘luster ages with respect to the models.," As a direct result of this, the KFF99 models, which have a tendency to systematically over-predict metallicities with respect to the models (see $\S$ \ref{subsec:TheLociofSWB-TypesontheSSPGrids}) ), systematically under-predict the cluster ages with respect to the models."
+ This is illustrated. in Figure 13.. where we compare the age and metallicity predictions of the? and. INET99 models using two dillerent metallicity indicators (iE). Mg») and the more age-sensitive 11.," This is illustrated in Figure \ref{fig:compare.ssp}, , where we compare the age and metallicity predictions of the and KFF99 models using two different metallicity indicators $\langle$ $\rangle$, $_2$ ) and the more age-sensitive $\beta$."
+ Phe clleet can be most clearly seen in the Lh? plane of the models., The effect can be most clearly seen in the $\langle$ $\rangle$ $\beta$ plane of the models.
+ The KEE99O-derivecl metallicities are systematically 0.20.5 dex higher than the ?--derivecl metallicities., The KFF99-derived metallicities are systematically 0.2–0.5 dex higher than the -derived metallicities.
+ This leads to cluster ages which are significantly vounger in the IXEE99 mocels (up to 6 Gar for old. ages) than those of2., This leads to cluster ages which are significantly younger in the KFF99 models (up to 6 Gyr for old ages) than those of.
+. However. surprisingly. the agreement between models for the ages of the voungest clusters is much better. even. though the agreement between their metallicities is poorest.," However, surprisingly, the agreement between models for the ages of the youngest clusters is much better, even though the agreement between their metallicities is poorest."
+ The origin of these dillerences are unclear. since both the mocdoels. use the same fitting-functions. but a possible explanation may fie in their adoption of dilferent input isochrones.," The origin of these differences are unclear, since both the models use the same fitting-functions, but a possible explanation may lie in their adoption of different input isochrones."
+ Of final. important note in Figure 13.. there is a clear olfset in metallicity between the (Fe}- and Mgo-derived metallicities. in the sense that the Ales index predicts metallicities 0.1 ~ 0.5 dex higher than ¢Fe}.," Of final, important note in Figure \ref{fig:compare.ssp}, there is a clear offset in metallicity between the $\langle$ $\rangle$ - and $_2$ -derived metallicities, in the sense that the $_2$ index predicts metallicities 0.1 $\sim$ 0.5 dex higher than $\langle$ $\rangle$."
+ However. we lind no evidence of a significant. systematic ollsct between our measured magnesium and iron indices.," However, we find no evidence of a significant systematic offset between our measured magnesium and iron indices."
+ As we showed in Section 4.. we were able to correct both these indices onto the Lick/IDS svstem.," As we showed in Section \ref{sec:TheSpectroscopicSystem}, we were able to correct both these indices onto the Lick/IDS system."
+ An alternative explanation is that the Fe? ancl Mg» indices do not track metallicity in the same manner: the Kel] measurements of the clusters are systematically lower than our Mg/LH] measurements., An alternative explanation is that the $\langle$ $\rangle$ and $_2$ indices do not track metallicity in the same manner; the [Fe/H] measurements of the clusters are systematically lower than our [Mg/H] measurements.
+ Such “a-enhancement” has been seen in the integrated spectra of elliptical galaxies (e.g. 1989: Gonzélez 1993: 2001: 2000a) ane recently in extragalactic globular clusters 2001. 2002).," Such $\alpha$ -enhancement"" has been seen in the integrated spectra of elliptical galaxies (e.g. 1989; $\acute{a}$ lez 1993; 2001; 2000a) and recently in extragalactic globular clusters 2001, 2002)."
+ Moreover. high-resolution spectroscopy of LAIC clusters giants suggests that. from Ο/Η] ratios. this is also the case for LMC clusters(2).," Moreover, high-resolution spectroscopy of LMC clusters giants suggests that, from [O/H] ratios, this is also the case for LMC clusters."
+. However. such an interpretation is complicated by the fact. that. at low metallicities. the differences between solar-caled and e-enhanced SSP mocdels are small little cbvnamic range. 2000).," However, such an interpretation is complicated by the fact that, at low metallicities, the differences between solar-scaled and $\alpha$ -enhanced SSP models are small little dynamic range, 2000)."
+ A detailed analysis of this important issue is bevond the scope of this paper. and we defer further discussion to future work.," A detailed analysis of this important issue is beyond the scope of this paper, and we defer further discussion to future work."
+ In ‘Table Bl. we list the age and metallicity predictions of the SSP models. using the Meg»-117 and £E5-LEz 47dices.," In Table B1, we list the age and metallicity predictions of the SSP models, using the $_2$ $\beta$ and $\langle$ $\rangle$ $\gamma_{\rm F}$."
+ To test these model predictions. we have collected age and metallicity estimates for the LAIC clusters. in our sample from the literature.," To test these model predictions, we have collected age and metallicity estimates for the LMC clusters in our sample from the literature."
+ These are also presented in Table Bl., These are also presented in Table B1.
+ These literature ages and metallicities come [roni a variety sources and are therefore. rather inhomogenous., These literature ages and metallicities come from a variety sources and are therefore rather inhomogenous.
+ Where possible. we have tried to minimise this inhomogencity whilst retaining a large enough sample for meaningful comparisons.," Where possible, we have tried to minimise this inhomogeneity whilst retaining a large enough sample for meaningful comparisons."
+ Cluster ages are preferentially taken from studies which have located the main sequence turn-olf in CMDs., Cluster ages are preferentially taken from studies which have located the main sequence turn-off in CMDs.
+ These have been supplemented with ages obtained by the giant-branch calibration of (1982)., These have been supplemented with ages obtained by the giant-branch calibration of (1982).
+ For the seven clusters withno previous spectroscopic- audor CALD-derivec age determinations. we have adopted a mean age which corresponds to their SWB type(sce Table 7)).," For the seven clusters withno previous spectroscopic- and/or CMD-derived age determinations, we have adopted a mean age which corresponds to their SWB type(see Table \ref{tab:swb.ages}) )."
+ We assign uncertainties by taking the 50 percentiles in their SWB age-range., We assign uncertainties by taking the 50 percentiles in their SWB age-range.
+ The majority ofthe literature cluster metallicities come from the ~ 2A rresolution Ca-triplet spectroscopy. of cluster τος giants by, The majority of the literature cluster metallicities come from the $\sim$ 2 resolution Ca-triplet spectroscopy of cluster red giants by
+there is a flux decrease at Obie 70.5. possibly due to the fact that part of the secondary star is shadowed by the accretion disk or to a warp in the disk itself.,"there is a flux decrease at $\phi_{binary}\sim$ 0.5, possibly due to the fact that part of the secondary star is shadowed by the accretion disk or to a warp in the disk itself."
+" ""he rise from the eclipse (Ó,;,,4 770.1-0.3) is steeper in UVWI (longer wavelength) than in UVW2.", The rise from the eclipse $\phi_{binary}\sim$ 0.1-0.3) is steeper in UVW1 (longer wavelength) than in UVW2.
+ Since. at these orbital phases. the accretion disk mainly contributes to the UV [lux. this is consistent with a scenario in which. after the eclipse. the regions of the disk farthest from the neutron star come into view first.," Since, at these orbital phases, the accretion disk mainly contributes to the UV flux, this is consistent with a scenario in which, after the eclipse, the regions of the disk farthest from the neutron star come into view first."
+ The spin period of Her N-1 has been found to vary from 1.23172 8 to 1.23782 s. with alternating phases of spin-up and spin-cown (Parmar et al.," The spin period of Her X-1 has been found to vary from $\sim$ 1.23772 s to 1.23782 s, with alternating phases of spin-up and spin-down (Parmar et al."
+ 1999. Oosterbrock at al.," 1999, Oosterbroek at al."
+ 2001)., 2001).
+ Furthermore. it can vary significantly. on relatively short time scales.," Furthermore, it can vary significantly on relatively short time scales."
+ The values of the spin period measured using the first three observations were reported by 1102: using the PN detector. we found 1.237774 s. 1.231753 hi and 1.237751 s. respectively.," The values of the spin period measured using the first three observations were reported by R02: using the PN detector, we found 1.237774 s, 1.237753 s and 1.237751 s, respectively."
+ Phe first measure confirmed the slow-down trend monitored byDeppoSax (Oosterbroek et al., The first measure confirmed the slow-down trend monitored by (Oosterbroek et al.
+ 2000) and Chandra (Burwitz. private communication): although we observed a slight decrease in the two following epochs.," 2000) and Chandra (Burwitz, private communication); although we observed a slight decrease in the two following epochs."
+ We performed. a search for pulsations using a Discrete Fourier Transform for all the new datasets. finding a strong evidence of pulsations onlv in one case (revolution. 415).," We performed a search for pulsations using a Discrete Fourier Transform for all the new datasets, finding a strong evidence of pulsations only in one case (revolution 415)."
+ This is the ον new exposure taken outside the low state. precisely during the main-on at $5;=0.02.," This is the only new exposure taken outside the low state, precisely during the main-on at $\Phi_{35}=0.02$."
+ The corresponding spin period. as obtained from the EPIC PN data. is L237777(1) s. where the number in parenthesis represents the error on the last digit as determined by using the Cash (1979). statistic.," The corresponding spin period, as obtained from the EPIC PN data, is 1.237777(1) s, where the number in parenthesis represents the error on the last digit as determined by using the Cash (1979) statistic."
+ The spin period is just. slightly increased with respect to the value measured. by in 2001., The spin period is just slightly increased with respect to the value measured by in 2001.
+ All the remaining datasets showed zumplitude spectra with only weak or no significant peaks close to the expected spin period., All the remaining datasets showed amplitude spectra with only weak or no significant peaks close to the expected spin period.
+ ln order to investigate further the possible presence of a modulation. we folded all the datasets on the most recent measured period. ie. 1.237774 s. by using 25 phase bins.," In order to investigate further the possible presence of a modulation, we folded all the datasets on the most recent measured period, i.e. 1.237777 s, by using 25 phase bins."
+ We then fitted a constant value to the resulting light curve and. determined thex7., We then fitted a constant value to the resulting light curve and determined the.
+ Vherefore. we expect a poorer fit in datasets with a stronger spin modulation.," Therefore, we expect a poorer fit in datasets with a stronger spin modulation."
+ Since the profile of the Folded light curve varies considerably in the soft ancl hard energy band. we split the events into two enerev ranges: 0.3-0.7 keV and 2-10 keV. The resulting is shown in Figure 3.. as a function of the 35 day phase.," Since the profile of the folded light curve varies considerably in the soft and hard energy band, we split the events into two energy ranges: 0.3-0.7 keV and 2-10 keV. The resulting is shown in Figure \ref{spin_var}, as a function of the 35 day phase."
+ For 1 degree of freedoms (the value of the period). the dillerence to the Git is Ay7=22.71 at the 90 percent conlidence level.," For 1 degree of freedom (the value of the period), the difference to the fit is $\Delta$ 2.71 at the 90 percent confidence level."
+ Using this method we find that. outside the main-on and short-on states there is evidence for significant modulation above 2 keV in the two datasets taken at Q4; 0.26 and 0.31.," Using this method we find that, outside the main-on and short-on states there is evidence for significant modulation above 2 keV in the two datasets taken at $\Phi_{35}$ =0.26 and 0.31."
+ In the first case (whieh is one of those already reported. by 102) the increases above the 90 percent confidence level if we reduce the number of phase bins from 25 to 20., In the first case (which is one of those already reported by R02) the increases above the 90 percent confidence level if we reduce the number of phase bins from 25 to 20.
+" Defining the amplitude as (max-min/moean) of a fitted sinusoid. we find an amplitude of 11 and 13 percent for (p5,—0.26 and 0.31. respectively."," Defining the amplitude as (max-min/mean) of a fitted sinusoid, we find an amplitude of 11 and 13 percent for $\Phi_{35}$ =0.26 and 0.31, respectively."
+ Previous observations of Ler X-1 during the low state include Coburn et al. (, Previous observations of Her X-1 during the low state include Coburn et al. (
+2000). who found a spin modulation of 13 percent (assuming the same cdelinition as ours) in the 3.18 keV range usingTE data in an anomalous low state. while Mihara ct al. (,"2000), who found a spin modulation of $\sim$ 13 percent (assuming the same definition as ours) in the 3–18 keV range using data in an anomalous low state, while Mihara et al. ("
+1991) determined an upper limit for the pulsed. fraction of 2.4 percent in the 1.2-37 keV energy range using data.,1991) determined an upper limit for the pulsed fraction of 2.4 percent in the 1.2-37 keV energy range using data.
+ In Figure 4.. we show the spin profiles of Her X-1 as derived [or those observations in which we detected. a significant modulation in the 0.3-0.7. keV or in the 2-10 keV energy band.," In Figure \ref{spin_hard}, we show the spin profiles of Her X-1 as derived for those observations in which we detected a significant modulation in the 0.3-0.7 keV or in the 2-10 keV energy band."
+ Due to the uncertainty in the spin period at each epoch. the spin phases are not on the same absolute scale.," Due to the uncertainty in the spin period at each epoch, the spin phases are not on the same absolute scale."
+ Therefore. only the relative phasing between the light curves in different energy. bands is physically meaningful (see 4.3)).," Therefore, only the relative phasing between the light curves in different energy bands is physically meaningful (see \ref{phase_shift}) )."
+ The change of the spin profile of Her. N-1 in cilferent energv bands has been well documented in the literature., The change of the spin profile of Her X-1 in different energy bands has been well documented in the literature.
+ In order to compare the spin profiles obtained by using with the previously published data. we take as reference. in the range above ~2 keV. the work of Decter et al. (," In order to compare the spin profiles obtained by using with the previously published data, we take as reference, in the range above $\sim 2$ keV, the work of Deeter et al. ("
+1998) and Scott et al. (,1998) and Scott et al. (
+2000).,2000).
+ These authors reported, These authors reported
+ ⊺⊺≋≺↸∖∙∶↴∙⊾∙∙⋀∖↕∏∑↸∖↥⋅∪∐↸∖↸∖↑⋜↕↕∙↕≝↭≺∖∖⊔⋜⋯≼,"TTS (e.g., Muzerolle et al. \cite{Mea98b}) )"
+⊔∐↑↕∐∖↕⋟↸∖↖↖↽∖⊽∫⇀⋀∖↕≼≓≽↴∖↴ for which οσοιId be estimated from veiling measurements (Muzerolle eta. 2003))., and in the few VLMOs for which could be estimated from veiling measurements (Muzerolle et al. \cite{Mea03}) ).
+ The underlying assuuption is that the eutire (Cas well as the other hydrogen mes) comes from accreting matter. and that the contanunation frou wind/jet cnussion is uceligible.," The underlying assumption is that the entire (as well as the other hydrogen lines) comes from accreting matter, and that the contamination from wind/jet emission is negligible."
+ This is not necessarily true in all stronely accreting TTS (sec. for example. Bacciotti et al. 2002)).," This is not necessarily true in all strongly accreting TTS (see, for example, Bacciotti et al. \cite{Bac02}) ),"
+ but is probably a good approximation for our objects (sec 83.1)., but is probably a good approximation for our objects (see 3.4).
+" For all objects we have adopted a stellar mass =550AL;.. radius 0.5 ad a maguctospleric truucation radius of 2.23R,."," For all objects we have adopted a stellar mass 50, radius 0.5 and a magnetospheric truncation radius of 2.2–3."
+. The results are not very seusitive to the exact values of these paralucters., The results are not very sensitive to the exact values of these parameters.
+ The fitting procedure gives for each object a value of The fits are of acceptable quality. with the exception of Cha Ha2. whose broad and flat profile (ccarly preseut iu all the observations) cannot be reproduced by the curreut nodels (similar profiles are observed in few other VEMOs and discussed in Muzerolle et al. 20033).," The fitting procedure gives for each object a value of The fits are of acceptable quality, with the exception of Cha 2, whose broad and flat profile (clearly present in all the observations) cannot be reproduced by the current models (similar profiles are observed in few other VLMOs and discussed in Muzerolle et al. \cite{Mea03}) )."
+ The fit can be somewhat iuproved by adopting au ad-hoc temperature xofile for he accretion flow. whereby the temperature has οσα reduced relative to the fiducial model iu the outer xwt of the flow near the disk.," The fit can be somewhat improved by adopting an ad-hoc temperature profile for the accretion flow, whereby the temperature has been reduced relative to the fiducial model in the outer part of the flow near the disk."
+ This results in less euidssion rear the line ceuter. the most uncertai1 part of the profiles eivoen the approxinatious used in the radiative transfer calculations (see Muzerolle et al. 20013).," This results in less emission near the line center, the most uncertain part of the profiles given the approximations used in the radiative transfer calculations (see Muzerolle et al. \cite{Mea01}) )."
+" Nevertheless. voth models for Cla IIo2 have and an acerction rate of the order of LO29Ίντι, which we will adopt in the following."," Nevertheless, both models for Cha 2 have and an accretion rate of the order of $10^{-10}$, which we will adopt in the following."
+ As discussed by Muzerolle et al. (20033).," As discussed by Muzerolle et al. \cite{Mea03}) ),"
+ the values of delerived from the pprofiles should be accurate wihin a factor of 35., the values of derived from the profiles should be accurate within a factor of $\sim 3-5$.
+" A second group of 3 objects. all in Cha EL. (Cha Hal. van‘ha Was. Cha Wah) has always narrow pprofiles, typical of cromeospleric activity."," A second group of 3 objects, all in Cha I, (Cha 1, Cha 3, Cha 5) has always narrow profiles, typical of cromospheric activity."
+ For these objects. we estimate au upper luit to oof about 101?/vr.. for which the magnetospheric aceretion inodels predict a. line ciuission below the detection luit.," For these objects, we estimate an upper limit to of about $10^{-12}$, for which the magnetospheric accretion models predict a line emission below the detection limit."
+ Iu Fig., In Fig.
+ 3 we compare our estimates of tto those of other VLMOSs and TTS from the literature., \ref{haw} we compare our estimates of to those of other VLMOs and TTS from the literature.
+ The fleure plots aas a function of the full width of, The figure plots as a function of the full width of
+A [fundamental property of the distribution of galaxies is clustering. manifested by the presence of groups and clusters of galaxies and quantitatively measured by the correlation,"A fundamental property of the distribution of galaxies is clustering, manifested by the presence of groups and clusters of galaxies and quantitatively measured by the correlation"
+totaling Raco Transients (RRS) are neutron stars which were discovered. only through their isolated. pulses (?)..,Rotating Radio Transients (RRATs) are neutron stars which were discovered only through their isolated pulses \citep{mll+06}.
+ Some. however. have later been detectable through »eriodicity searches.," Some, however, have later been detectable through periodicity searches."
+ Phe average intervals between detected »ulses range from a few minutes to a few hours and. pulses rave durations between 2 and 30 ms., The average intervals between detected pulses range from a few minutes to a few hours and pulses have durations between 2 and 30 ms.
+ Thus far. 750 RATS iive been identified (2772?77).. including the original 11 rom ?..," Thus far, $\sim$ 50 RRATs have been identified \citep[][]{hrk+08,dcm+09,klk+09,bb10,kkl+11,bbj+11}, including the original 11 from \citet{mll+06}."
+ Periods ranging from 0.1 to 8 seconds have ⋡⊔⊔⊔↓⊓⇤⊔↓⊔∏∪↓−⊽⋟ thes sources., Periods ranging from 0.1 to 8 seconds have been measured for 29 of these sources.
+ dirivatives rave been d forol 14. allow inference.1eriod of spindown properties meatusuch as characteristicingages and surface dipole magnetic fields (???)..," Period derivatives have been measured for 14, allowing inference of spin--down properties such as characteristicages and surface dipole magnetic fields \citep[][]{mlk+09,lmk09,kkl+11}."
+ The periods. ancl magnetic ields of IRIUGVIS are larger than those of normal pulsars. out. the distributions of other spindown properties such as spindown energv loss rate and characteristic age are similar (?7)..," The periods and magnetic fields of RRATs are larger than those of normal pulsars, but the distributions of other spin–down properties such as spin–down energy loss rate and characteristic age are similar \citep{mlk+09}."
+ Despite this overall trend. the properties of individual RRATS vary considerably.," Despite this overall trend, the properties of individual RRATs vary considerably."
+ Four RRATS. including PSRs 1419 ancl 11911330. have spindown properties consistent with the bulk of the normal radio pulsar population and two others. PSRs 5759 and 6026. have properties similar to normal. older pulsars.," Four RRATs, including PSRs $-$ 1419 and J1913+1330, have spin--down properties consistent with the bulk of the normal radio pulsar population and two others, PSRs $-$ 5759 and$-$ 6026, have properties similar to normal, older pulsars."
+ Four others. PSRs 4406. 4417. JLSOT 2557. and 1419. lie just above the. racio ‘death[ine (e.g.2?7).," Four others, PSRs $-$ 4406, $-$ 4417, $-$ 2557, and $-$ 1419, lie just above the radio `death–line' \citep[e.g.][]{cr93,zgd07}."
+. However. some have more unusual spindown properties.," However, some have more unusual spin–down properties."
+ PSRs 4816. 91540 0257. and JISS4| 0306.tΞ lie in Ean empty region⊓ ofP?P space between the normal radio pulsars and. isolated neutron stars (NINS) and PSR L458 has a high magnetic field of 5.1047 C1. Because of the dillieulties in detecting these sporaclic objects. the total Galactic population of RRATSs likely outnumbers that of normal radio pulsars (?).. though it is possible that the populations are evolutionarily related (?)..," PSRs $-$ 4316, $-$ 0257, and $+$ 0306, lie in an empty region of $P-\dot{P}$ space between the normal radio pulsars and isolated neutron stars (XINS) and PSR $-$ 1458 has a high magnetic field of $5\times10^{13}$ G. Because of the difficulties in detecting these sporadic objects, the total Galactic population of RRATs likely outnumbers that of normal radio pulsars \citep{mll+06}, , though it is possible that the populations are evolutionarily related \citep{kk+08}. ."
+ Several ideas have heen presented. about the nature, Several ideas have been presented about the nature
+"whereJ, is the first-order Bessel [unction and where is (he projected baseline at the stars position. 0p is (he apparent UD angular cliameter of the star. and A is the wavelength of the observation.","where$_1$ is the first-order Bessel function and where is the projected baseline at the star's position, $\theta_{\rm UD}$ is the apparent UD angular diameter of the star, and $\lambda$ is the wavelength of the observation."
+" The limb-darkened (LD) relationship incorporating the linearlimb darkening coefficient. jp. 1974) is givenbv: These fitsresulted in Oy) = 0.366+0.024 mas and 0,4, = 0.38772:0.024 mas. the latter incorporating fy = 0.36 taken from Claretοἱal.(1995) after adopting log g = 4.5 and Τι.ell = 5000 K for HD 189733 (see Fieurei 4))."," The limb-darkened (LD) relationship incorporating the linearlimb darkening coefficient $\mu_{\lambda}$ \citep{1974MNRAS.167..475B} is givenby: These fitsresulted in $\Theta_{\rm UD}$ = $0.366 \pm 0.024$ mas and $\Theta_{\rm LD}$ = $0.377 \pm 0.024$ mas, the latter incorporating $\mu_{\lambda}$ = 0.36 taken from \citet{1995A&AS..114..247C} after adopting log $g$ = 4.5 and $T_{\rm eff}$ = 5000 K for HD 189733 (see Figure \ref{HD189733_viscurve}) )."
+ The reduced 4? minimization in both cases vielded a value of 1.593. and the errors quotedare for an increase of the 47 value of 1.0. that is. the confidence interval.," The reduced $\chi^2$ minimization in both cases yielded a value of 1.593, and the errors quotedare for an increase of the $\chi^2$ value of 1.0, that is, the confidence interval."
+ Dividing this V7 by the number of degrees of freedom. which in our case is 8. vields 0.199. which is much less than 1.0 showing that the fit is quite good and (hat our error estimates for (he visibility points are conservative.," Dividing this $\chi^2$ by the number of degrees of freedom, which in our case is 8, yields 0.199, which is much less than 1.0 showing that the fit is quite good and that our error estimates for the visibility points are conservative."
+ If we rescale these errors bars to force \7 to be equal to the number of degrees of freedom. which assumes (hal (here are no svsteniatics in the measurements. (μον are approximately half the size as thev are shown in Figure 2 and would also reduce our final error estimates by a [actor of (wo.," If we rescale these errors bars to force $\chi^2$ to be equal to the number of degrees of freedom, which assumes that there are no systematics in the measurements, they are approximately half the size as they are shown in Figure 2 and would also reduce our final error estimates by a factor of two."
+ llowever. we will remain conservative and cont(ünue to use the error estimate based on the raw AC> value.," However, we will remain conservative and continue to use the error estimate based on the raw $\chi^2$ value."
+ Anpriori estimate of the angular size of IID 189733 is a parameter of considerable interest. because (he size of WD 189733b is determined only relative to the size of ils parent star Irom (he photometric transit Gming data.," An estimate of the angular size of HD 189733 is a parameter of considerable interest, because the size of HD 189733b is determined only relative to the size of its parent star from the photometric transit timing data."
+ Bakosetal.(2006b) consider no less than four separate methods in their investigation of the svstem: V—A color angular radius prediction (Ixervellaetal.2004).. temperature radius. isochrone radii from Girardietal.(2002) ad Daralfeetal.(1998).. and the Johnson V - 2MAÀSS They calibration of Masanaοἱal. (2006)..," \citet{2006ApJ...650.1160B} consider no less than four separate methods in their investigation of the system: $V-K$ color angular radius prediction \citep{2004A&A...426..297K}, temperature radius, isochrone radii from \citet{2002A&A...391..195G} and \citet{1998A&A...337..403B}, and the Johnson V - 2MASS $T_{\rm EFF}$ calibration of \citet{2006A&A...450..735M}. ."
+ None of (hese approaches appears to have much merit. since the only primary cata we have been able to find in the literature were Tvcho By. Vr (Bessell 2000).. Strómmngren ubiy (Olsen 1993).. and 221ASS JN photometry (Cutriοἱal. 2003)... ," None of these approaches appears to have much merit, since the only primary data we have been able to find in the literature were Tycho $B_T$ , $V_T$ \citep{2000PASP..112..961B}, , Strömmgren \citep{1993A&AS..102...89O}, , and 2MASS $JHK$ photometry \citep{2003tmc..book.....C}. ."
+No spectroscopy or measures of log y appear to be available in the literature. nor dodirect measures of Johnson," No spectroscopy or measures of log $g$ appear to be available in the literature, nor dodirect measures of Johnson"
+magnetograms (dotted line) Naf-oquy~) where the numerator and d=1.6 Mm are found from fitting.,"magnetograms (dotted line) N_n(z), where the numerator and $d=1.6$ Mm are found from fitting."
+ The null column above a given height does vary between magnetograms. as shown in relfig:smrv..," The null column above a given height does vary between magnetograms, as shown in \\ref{fig:smry}."
+ There appear to be random fluctuations superimposed on a much slower variation., There appear to be random fluctuations superimposed on a much slower variation.
+ We estimated the dispersion about the trend. 2.7x10.!Mim.? for the one-minute niagnetograms and 2.5x10!Mm7 for five-minute averages. by subtracting a 90-day running mean.," We estimated the dispersion about the trend, $2.7\times10^{-4}\,{\rm Mm}^{-2}$ for the one-minute magnetograms and $2.5\times10^{-4}\,{\rm Mm}^{-2}$ for five-minute averages, by subtracting a 90-day running mean."
+ Considering the entire data set of one-minute magnetograms. including its slow trend. we lind the column of null points above z=1.5 Mm to be (The errors quoted above represent rms variation: (hie error in the mean is far smaller.)," Considering the entire data set of one-minute magnetograms, including its slow trend, we find the column of null points above $z=1.5$ Mm to be (The errors quoted above represent rms variation; the error in the mean is far smaller.)"
+ The value from five-minute magnetogranis is slightly lower: (2.9220.3)xLO7., The value from five-minute magnetograms is slightly lower: $(2.9\pm 0.3)\times10^{-3}$.
+ We believe (his value is lower partly because of motional blurring during (the averaging interval. so we adopt the oneaninuteanagnetogram value as (he more accurate of (he two.," We believe this value is lower partly because of motional blurring during the averaging interval, so we adopt the one-minute-magnetogram value as the more accurate of the two."
+" We believe the slow variation in ;V,(1.5) is due to gradual changes in the NTF of ihe imagine svstem.", We believe the slow variation in $N_n(1.5)$ is due to gradual changes in the MTF of the imaging system.
+ The reported focus position of the instrument is shown as a solid curve long the bottom., The reported focus position of the instrument is shown as a solid curve long the bottom.
+ The locus is periodically changed to compensate for changes in (he instrument itself., The focus is periodically changed to compensate for changes in the instrument itself.
+ Our estimate of the MTE was performed (through cross-comparison to the high resolution images. plotted as squares. all while the instrument was in focus position 5.," Our estimate of the MTF was performed through cross-comparison to the high resolution images, plotted as squares, all while the instrument was in focus position 5."
+ This is the MTE by which we are correcting the spectrum. so we believe that the value in the middle of this interval is probably the most accurate.," This is the MTF by which we are correcting the spectrum, so we believe that the value in the middle of this interval is probably the most accurate."
+ The ΣΤΕ correction raises the low resolution columns by about40%.. thereby making them more consistent with high resolution values.," The MTF correction raises the low resolution columns by about, thereby making them more consistent with high resolution values."
+" Omitting that step gives null columns which could be considered lower bounds on the actual values: meclians. are IN,(1.5)-x=2.210Mmu>στ forH one-minute. and 2.]x10""AimM>7 forH five-minute."," Omitting that step gives null columns which could be considered lower bounds on the actual values: medians are $N_n(1.5)=2.2\times10^{-3}\,{\rm Mm}^{-2}$ for one-minute and $2.1\times10^{-3}\,{\rm Mm}^{-2}$ for five-minute."
+H. A comparable plot from the previous solar minimum.," A comparable plot from the previous solar minimum,"
+"define fmin and fmaz as the minimum and maximum value of f in the region R and a value fo such that One has then that for each value of f€(fiin,fmaz) there exists a value of a€(oj,o) such that","define $f_{min}$ and $f_{max}$ as the minimum and maximum value of $f$ in the region $\cal{R}$ and a value $f_0$ such that One has then that for each value of $f \in (f_{min}, f_{max})$ there exists a value of $\alpha \in (\alpha_l, \alpha_u)$ such that"
+"lifetime of the RG, and is given by If we define 6M to be the typical amount of mass lost in the collision, then the mass loss rate is To calculate an upper limit for the contribution of RG - MS star collisions to the mass loss rate, we assume all RG and MS stars have masses of 1Mo, and that the entire RG is destroyed in the collision.","lifetime of the RG, and is given by If we define $\delta M$ to be the typical amount of mass lost in the collision, then the mass loss rate is To calculate an upper limit for the contribution of RG - MS star collisions to the mass loss rate, we assume all RG and MS stars have masses of $1\mathrm{ M_{\odot}}$, and that the entire RG is destroyed in the collision."
+" Collisions involving 1Mo RGs yield an upper limit, because there is not an appreciable amount of RGs with masses less than »1Mo due to their MS lifetimes being greater than the age of the Galaxy."," Collisions involving $1\mathrm{M_{\odot}}$ RGs yield an upper limit, because there is not an appreciable amount of RGs with masses less than $\sim 1\mathrm{M_{\odot}}$ due to their MS lifetimes being greater than the age of the Galaxy."
+" For RGs with masses greater than 1M, the amount they contribute to the mass loss rate is a competition between their lifetimes and radii."," For RGs with masses greater than $1 \mathrm{M_{\odot}}$, the amount they contribute to the mass loss rate is a competition between their lifetimes and radii."
+ Red giant lifetimes decrease with mass (thereby decreasing the time they have to collide) and their radii increase with mass (thereby increasing the cross section)., Red giant lifetimes decrease with mass (thereby decreasing the time they have to collide) and their radii increase with mass (thereby increasing the cross section).
+ In their Fig., In their Fig.
+" 3, Daleetal.(2009) clearly show that the number of collisions decreases with increasing RG mass, indicating that the brevity of their lifetime wins over their large sizes."," 3, \citet{dale:2009} clearly show that the number of collisions decreases with increasing RG mass, indicating that the brevity of their lifetime wins over their large sizes."
+" One solar mass MS impactors should yield approximately an upper limit to the mass loss rate, since 1Meo MS stars are the most common for the PDMFs under consideration."," One solar mass MS impactors should yield approximately an upper limit to the mass loss rate, since $\sim 1\mathrm{M_{\odot}}$ MS stars are the most common for the PDMFs under consideration."
+ Since we assume that the entire RG is destroyed in the collision 6M= 1Mo., Since we assume that the entire RG is destroyed in the collision $\delta M = 1\mathrm{M_{\odot}}$ .
+" For the case that all impactors are Μο MS stars, we calculate Ίιπα(Τραι) from equation (38)) by noting that n.(rgai)=p«(rgai)/(1Mc)."," For the case that all impactors are $1\mathrm{M_{\odot}}$ MS stars, we calculate $n_{RG} (r_{gal})$ from equation \ref{eqn:n_rg}) ) by noting that $n_{\star}(r_{gal}) = \rho_{\star}(r_{gal})/(1M_{\odot})$."
+" For self-consistency, we must truncate P(rgai) at 1 for all P(rgai)>1 since we are considering the case where one collision destroys the entire star."," For self-consistency, we must truncate $P(r_{gal})$ at 1 for all $P(r_{gal}) >1$ since we are considering the case where one collision destroys the entire star."
+ We plot equation (40)) for this calculation in Fig. 16.., We plot equation \ref{eqn:mdot_rg}) ) for this calculation in Fig. \ref{fig:red_giant_cont}.
+ The discontinuity is due to our truncating P(rgai) at 1., The discontinuity is due to our truncating $P(r_{gal})$ at 1.
+" The figure shows that the mass loss rate for RG-MS star collisions never exceeds 10-°Moyr7', well below typical dM/dlnrgat for values for MS - MS collisions (see Figs."," The figure shows that the mass loss rate for RG-MS star collisions never exceeds $10^{-5}\mathrm{M_{\odot}yr^{-1}}$ , well below typical $d\dot M/dlnr_{gal}$ for values for MS - MS collisions (see Figs."
+ 9 and 10))., \ref{fig:dmdot_dlnr_v_r} and \ref{fig:dir_indir}) ).
+" Moreover, in their hydrodynamic simulations, Daleetal.(2009) note that in a typical RG - MS star collision, at most ~10% of the RG envelope is lost to the RG."," Moreover, in their hydrodynamic simulations, \citet{dale:2009} note that in a typical RG - MS star collision, at most $\sim 10\%$ of the RG envelope is lost to the RG."
+ We therefore conclude that the contribution of RGs to the total mass loss rate in the central parsec of the Galaxy is negligible., We therefore conclude that the contribution of RGs to the total mass loss rate in the central parsec of the Galaxy is negligible.
+" The figure shows that by rgo;= 0.06pc, the mass loss rate for RG-MS star collisions is at most about 10-°Moyr7!."," The figure shows that by $r_{gal} = 0.06$ pc, the mass loss rate for RG-MS star collisions is at most about $10^{-6}\mathrm{M_{\odot}yr^{-1}}$."
+" It is thus possible that for MS-MS collisions, values of Mi» and « that results in total mass loss rates just below 107Moyr-! could be pushed past this threshold with the addition of mass loss due to RG collisions."," It is thus possible that for MS-MS collisions, values of $M_{min}$ and $\alpha$ that results in total mass loss rates just below $10^{-5}\mathrm{M_{\odot}yr^{-1}}$ could be pushed past this threshold with the addition of mass loss due to RG collisions."
+" However, we believe that this is unlikely for two reasons."," However, we believe that this is unlikely for two reasons."
+" The inclusion of the factor, ¢, when calculating the total mass loss rate (see equation (26))) will reduce the mass loss by at least a factor of 0.6 (see Fig. 12))."," The inclusion of the factor, $\zeta$, when calculating the total mass loss rate (see equation \ref{eq:total_mass_loss_rate}) )) will reduce the mass loss by at least a factor of 0.6 (see Fig. \ref{fig:zeta}) )."
+" Also, as noted by the hydrodynamic simulations of Daleetal.(2009),, for a typical RG - MS star collision, at most ~10% of the RG envelope is lost to the RG."," Also, as noted by the hydrodynamic simulations of \citet{dale:2009}, for a typical RG - MS star collision, at most $\sim 10\%$ of the RG envelope is lost to the RG."
+ This will reduce dM/dlnrga; for RG-MS collisions by another factor of 10., This will reduce $d \dot M /dlnr_{gal}$ for RG-MS collisions by another factor of 10.
+" We have have derived novel, analytical methods for calculating the amount of mass loss from indirect and direct stellar collisions in the Galactic centre."," We have have derived novel, analytical methods for calculating the amount of mass loss from indirect and direct stellar collisions in the Galactic centre."
+" Our methods compares very well to hydrodynamic simulations, and do not require costly amounts of computation time."," Our methods compares very well to hydrodynamic simulations, and do not require costly amounts of computation time."
+ We have also computed the total mass loss rate in the Galactic centre due to stellar collisions., We have also computed the total mass loss rate in the Galactic centre due to stellar collisions.
+" Mass loss from direct collisions dominates at Galactic radii below ~0.1pc, and thereafter indirect collisions dominate the total mass loss rate."," Mass loss from direct collisions dominates at Galactic radii below $\sim0.1\mathrm{pc}$, and thereafter indirect collisions dominate the total mass loss rate."
+" Since the amount of stellar material lost in the collision depends upon the masses of the colliding stars, the total mass loss rate depends upon the PDMF."," Since the amount of stellar material lost in the collision depends upon the masses of the colliding stars, the total mass loss rate depends upon the PDMF."
+" We find that the calculated mass loss rate is sensitive to the PDMF used, and can therefore be used to constrain the PDMF in the Galactic centre."," We find that the calculated mass loss rate is sensitive to the PDMF used, and can therefore be used to constrain the PDMF in the Galactic centre."
+" As summarized by Fig. 13,,"," As summarized by Fig. \ref{fig:param_space},"
+ our calculations rule out aS;1.25 and Mmin27M in the Mj;—o parameter space., our calculations rule out $\alpha \lesssim 1.25$ and $M_{\min} \gtrsim 7 \mathrm{M_{\odot}}$ in the $M_{min}-\alpha$ parameter space.
+" Finally, we have used our constraints on the PDMF in the Galactic centre to constrain the IMF to have a power-law slope 2 0.4to 0.9 depending on the star formation history of the Galactic centre."," Finally, we have used our constraints on the PDMF in the Galactic centre to constrain the IMF to have a power-law slope $\gtrsim $ 0.4to 0.9 depending on the star formation history of the Galactic centre."
+" 'This work was supported in part by the National Science Foundation Graduate Research Fellowship, NSF grant AST-0907890 andNASA grants NNX08ALA43G and NNA09DB30A."," This work was supported in part by the National Science Foundation Graduate Research Fellowship, NSF grant AST-0907890 andNASA grants NNX08AL43G and NNA09DB30A."
+we can easily deduce [rom (3.5)) that Wi€W3 Since WircWe. we cau apply inequality (1.lo ) to the funetion Wi. aud. having (3.7)) iu αλά». the proof of Proposition 3.1. directly follows if we can show that auc this will be cloue in the lortheoming arguments.,"we can easily deduce from \ref{5s5}) ) that $\Psi \tilde{u}\in
+W_{2}^{2,1}(\R^{2})$, and Since $\Psi \tilde{u}\in W_{2}^{2,1}(\R^{2})$, we can apply inequality \ref{cara_eq2}) ) to the function $\Psi \tilde{u}$, and, having \ref{5s5_1}) ) in hands, the proof of Proposition \ref{prop1} directly follows if we can show that and this will be done in the forthcoming arguments."
+" In all what follows. it will be useful to deal with an equivalent norm of the BALO, space."," In all what follows, it will be useful to deal with an equivalent norm of the $BMO_p$ space."
+ This norm is given by the following lemma., This norm is given by the following lemma.
+ The proof of this lemuna is direct., The proof of this lemma is direct.
+ It suffices to see that for auy c€R. we have: which imiuediately. gives: hence aud the equivalence of the two norms follows.," It suffices to see that for any $c\in \R$, we have: which immediately gives: hence and the equivalence of the two norms follows."
+ α From uow on. aud for the sake of simplicity. we will denote: The following lemma gives an estimate of iuf/[u—e| on sinall parabolic cubes.," $\hfill{\blacksquare}$ From now on, and for the sake of simplicity, we will denote: The following lemma gives an estimate of $\displaystyle \inf_{c\in
+ \R}\inm_{Q}|u-c|$ on small parabolic cubes."
+Perna 2001).,Perna 2001).
+ Another interesting property hintecl by the GIAWOO sample of superdisks was that the hot spots in the radio lobe pair are usually located more svmametrically about the superdisk's mid-plane than about the host galaxy. possibly signifving the galaxys motion during the [CIN phase (see 44).," Another interesting property hinted by the GKW00 sample of superdisks was that the hot spots in the radio lobe pair are usually located more symmetrically about the superdisk's mid-plane than about the host galaxy, possibly signifying the galaxy's motion during the AGN phase (see 4)."
+ In high recdshift racio galaxies. the apparent asvmmetry of clilfuse Lya emission. which is a sensitive tracer of dust. can be understood if the brighter Lvo emission is associated with the radio lobe on the near side of the nucleus and thus not obscured. by any. dust. present. in the superdisk (GINNOQ).," In high redshift radio galaxies, the apparent asymmetry of diffuse $\alpha$ emission, which is a sensitive tracer of dust, can be understood if the brighter $\alpha$ emission is associated with the radio lobe on the near side of the nucleus and thus not obscured by any dust present in the superdisk (GKW00)."
+ Because the sharp-edeed morphology would only be noticed when the radio jets are quite close to the plane of the sky. they are subject to a strong negative selection elfect. and. CGINWOO argued that even though the observed cases of superdisk are few. the phenomenon max not be rare.," Because the sharp-edged morphology would only be noticed when the radio jets are quite close to the plane of the sky, they are subject to a strong negative selection effect, and GKW00 argued that even though the observed cases of superdisk are few, the phenomenon may not be rare."
+ Initially we proposed that the superdisk is primarily mace of the interstellar medium bound to the radio galaxy (RO) itself. perhaps originating from the gas belonging to gaserich disk galaxies previously captured hy the giant elliptical host of the powerful RG (e.g. Statler MleNamara 2002).," Initially we proposed that the superdisk is primarily made of the interstellar medium bound to the radio galaxy (RG) itself, perhaps originating from the gas belonging to gas-rich disk galaxies previously captured by the giant elliptical host of the powerful RG (e.g., Statler McNamara 2002)."
+ We argued that the tidal stretching anc heating of that gas during the capture was sullicient to produce very large fat pancakes (CGopal-Ixrishna Nath 1997: GINWOO)., We argued that the tidal stretching and heating of that gas during the capture was sufficient to produce very large fat pancakes (Gopal-Krishna Nath 1997; GKW00).
+" Since in some extreme cases (c.g. 476: Saripalli. Subrahmanvan Ucava Shankar 2002) the width of the superdisk was found to run into several hundred kiloparsecs. an alternative possibility was also put forward in GIN00. according to which at least some supercisks trace the gaseous filaments of the “cosmic web""."," Since in some extreme cases (e.g., $-$ 476; Saripalli, Subrahmanyan Udaya Shankar 2002) the width of the superdisk was found to run into several hundred kiloparsecs, an alternative possibility was also put forward in GKW00, according to which at least some superdisks trace the gaseous filaments of the “cosmic web""."
+" Earlier. noticing the straight. abrupt inner boundaries of the lobes in a raclio ealaxy (3€ 227). Black et ((1992) had inferred a ""disc of central dense. cold gas with axis coincident with that of he source."," Earlier, noticing the straight, abrupt inner boundaries of the lobes in a radio galaxy (3C 227), Black et (1992) had inferred a “disc of central dense, cold gas with axis coincident with that of the source""."
+ Very recently. Gergely Bicrmann (2007) have ooposed that superdisks can be carved out by the powerful wind produced due to à rapid jet precession during the impencding merger of two supermassive black holes belonging ο à pair of merging galaxies.," Very recently, Gergely Biermann (2007) have proposed that superdisks can be carved out by the powerful wind produced due to a rapid jet precession during the impending merger of two supermassive black holes belonging to a pair of merging galaxies."
+ A popular explanation for the emission gaps. in eeneral. invokes a squeezing or pinching of the (lighter) svnehrotron asma of the radio bridge in the middle. by the denser. ueher pressure ISM. or circum-galactic medium. associated with the parent galaxy (e.g.. Seheuer 1974: Croston ct 22004b).," A popular explanation for the emission gaps, in general, invokes a squeezing or pinching of the (lighter) synchrotron plasma of the radio bridge in the middle, by the denser, higher pressure ISM, or circum-galactic medium, associated with the parent galaxy (e.g., Scheuer 1974; Croston et 2004b)."
+ While it is presently. unclear how this process can account for the gaps with sharp. straight. boundaries. the existence of such a medium is indeed. supported by. the detection. of resolved. X-ray emission. from many massive ellipticals hosting double radio sources at small to medium redshifts ΕΠardeastle Worrall 2000).," While it is presently unclear how this process can account for the gaps with sharp, straight boundaries, the existence of such a medium is indeed supported by the detection of resolved X-ray emission from many massive ellipticals hosting double radio sources at small to medium redshifts Hardcastle Worrall 2000)."
+" Vhis ISM is likely to be built up through a gradual accumulation of the gas lost by the stellar population over the Hubble time and heated to ~10"" 101 Ix due to mixing within the galaxy dominated by random motions (Drishenti Mathews 1907 and references therein: OSullivan. Forbes Ponman 2001)."," This ISM is likely to be built up through a gradual accumulation of the gas lost by the stellar population over the Hubble time and heated to $\sim 10^6$ – $10^7$ K due to mixing within the galaxy dominated by random motions (Brighenti Mathews 1997 and references therein; O'Sullivan, Forbes Ponman 2001)."
+ On the other hand. massive cllipticals at high redshifts are expected to lack a significant amount of high-pressure ISM.," On the other hand, massive ellipticals at high redshifts are expected to lack a significant amount of high-pressure ISM."
+ Then the medium. arounc even massive galaxies situated in groups ancl clusters will resemble. peaks in the interealactic medium (GAL) (ancl not the higher pressure intracluster medium which is typical of their lower redshift counterparts) due to an early stage of eas aceretion from the cosmic web and its virialization., Then the medium around even massive galaxies situated in groups and clusters will resemble peaks in the intergalactic medium (IGM) (and not the higher pressure intracluster medium which is typical of their lower redshift counterparts) due to an early stage of gas accretion from the cosmic web and its virialization.
+ Basecl on. sensitive optical/X-ray observations ancl cosmological simulations. clusters at z>1.5 are believed το be. proto-clusters which would have grown into the present-day clusters. by the process of virialization ancl merger with other clusters or galaxy groups (ασ. Miley ct 22004: Ford οἱ 22004: also. Rowley. Thomas lav 2004: Motl et 22004: Tormen. Moscardini Yoshicla 2004).," Based on sensitive optical/X-ray observations and cosmological simulations, clusters at $z >~ 1.5 $ are believed to be proto-clusters which would have grown into the present-day clusters by the process of virialization and merger with other clusters or galaxy groups (e.g., Miley et 2004; Ford et 2004; also, Rowley, Thomas Kay 2004; Motl et 2004; Tormen, Moscardini Yoshida 2004)."
+ Significant concomitant enrichment of the ICM. is also expected: via galactic winds and ram-pressure stripping of the galaxy »opulation (e.g.. Ixapferer et 22007).," Significant concomitant enrichment of the ICM is also expected via galactic winds and ram-pressure stripping of the galaxy population (e.g., Kapferer et 2007)."
+ Our premise that distant radio galaxies have a circum-galactic mecium at a significantly: lower cmpcrature/pressure than. co nearer ones is supported w the deep observations which show that speetroscopically selected: high redshift (2~1) clusters ancl groups of galaxies are under-Iuminous. compared to heir low redshift counterparts (Fang et 22006).," Our premise that distant radio galaxies have a circum-galactic medium at a significantly lower temperature/pressure than do nearer ones is supported by the deep observations which show that spectroscopically selected high redshift $z \sim 
+1$ ) clusters and groups of galaxies are under-luminous, compared to their low redshift counterparts (Fang et 2006)."
+ A similar result is reported for medium. recshilt groups of galaxies (Spiegel. Paerels Scharf 2007).," A similar result is reported for medium redshift groups of galaxies (Spiegel, Paerels Scharf 2007)."
+ Ehe lower X-ray output. probably indicates that these high-2 systems have vet to acerete adequate hot gas and/or being dynamically vounger svstems. they are not vet. virlalized.," The lower X-ray output probably indicates that these $z$ systems have yet to accrete adequate hot gas and/or being dynamically younger systems, they are not yet virialized."
+ Lt is thus expected that the (forming) clusters at high redshifts would be lacking a dense high-pressure medium. found: within nearby clusters and. groups of galaxies., It is thus expected that the (forming) clusters at high redshifts would be lacking a dense high-pressure medium found within nearby clusters and groups of galaxies.
+ Lt may be noted that although N-rav emission. extended on at. least 100 kpe scales has been detected around the 2=1.286 quasar 3C294 and the z=2.48 I 4€ 23.56 from deep. observations. it is more likely to. be inverse Compton boostecdl cosmic microwave background. as its energy density is much. greater at such high. redshifts (Fabian οἱ 22003: Johnson et 22007: see also. Celotti Fabian 2004).," It may be noted that although X-ray emission extended on at least 100 kpc scales has been detected around the $z= 1.786$ quasar 3C294 and the $z=2.48$ RG 4C 23.56 from deep observations, it is more likely to be inverse Compton boosted cosmic microwave background, as its energy density is much greater at such high redshifts (Fabian et 2003; Johnson et 2007; see also, Celotti Fabian 2004)."
+ These considerations motivate us to explore in this (uper an alternative mechanism that could. be ellective in 16 formation of emission gaps between the radio lobes in ügher recshift racio galaxies., These considerations motivate us to explore in this paper an alternative mechanism that could be effective in the formation of emission gaps between the radio lobes in higher redshift radio galaxies.
+ Non-relativistic winds. with speeds often exceeding LO? kim  and a mass outllow rates often greater than 1M. | (eg Soker Pizzolato 2005: Brighenti Mathews 2006: Crenshaw Ixraemer 2007: Tremonti. Moustakas Diamond-Stanic 2007) are now thought to be an integral xwt of the active galactic nucleus (AGN) phenomenon. and can be launched from accretion disks via several mechanisms (c.g... Naravan Yi 1994: Wonniel Ixartje 1994: Blandford Degelman 1990: Das et 22001).," Non-relativistic winds, with speeds often exceeding $^3$ km $^{-1}$ and a mass outflow rates often greater than 1 $_{\odot}$ $^{-1}$ (e.g., Soker Pizzolato 2005; Brighenti Mathews 2006; Crenshaw Kraemer 2007; Tremonti, Moustakas Diamond-Stanic 2007) are now thought to be an integral part of the active galactic nucleus (AGN) phenomenon, and can be launched from accretion disks via several mechanisms (e.g., Narayan Yi 1994; Könnigl Kartje 1994; Blandford Begelman 1999; Das et 2001)."
+ Evidence for such outllows. comes from their absorption of the underlving continuum in the UV and/or X-ravs (e.g... PDS 456. Reeves. OBrien Ward 2003: PC 1211] 143. Pounds ct 22003: NGC 1097. Storchi-Dergmann ct 2003: see Crenshaw. Kraemer (ποσο 2003 for a review).," Evidence for such outflows comes from their absorption of the underlying continuum in the UV and/or X-rays (e.g., PDS 456, Reeves, O'Brien Ward 2003; PG $+$ 143, Pounds et 2003; NGC 1097, Storchi-Bergmann et 2003; see Crenshaw, Kraemer George 2003 for a review)."
+ Such relatively slow winds and relativistic jets may coexist in many objects BBinney 2004: Cirege. Decker de Vries 2006).," Such relatively slow winds and relativistic jets may coexist in many objects Binney 2004; Gregg, Becker de Vries 2006)."
+ As these non-relativistic winds. together with any existing relativistic jets. are believed: to carry a mechanical luminosity far in excess of the svnchrotron luminosity (e.g... Willott et 11999). they are. capable of significantly. influencing important phenomena such as cooling Lows in clusters (e... Omma Binney 2004: Soker Pizzolato 2005: Brighenti Alathews 2006). ancl even the structure formation in the universe (e.g. Rawlings 2003: (οσο et 2006).," As these non-relativistic winds, together with any existing relativistic jets, are believed to carry a mechanical luminosity far in excess of the synchrotron luminosity (e.g., Willott et 1999), they are capable of significantly influencing important phenomena such as cooling flows in clusters (e.g., Omma Binney 2004; Soker Pizzolato 2005; Brighenti Mathews 2006), and even the structure formation in the universe (e.g., Rawlings 2003; Gregg et 2006)."
+ Lt is certainly possible that. the jets contain, It is certainly possible that the jets contain
+The 1.4 GHz VLA image of ? shows a faint lobe south of the core. with a weak “hot spot” in PA=160°.,"The 1.4 GHz VLA image of \citet{Rector03} shows a faint lobe south of the core, with a weak “hot spot” in $\degr$."
+ The 5 Gllz VLBA map however shows a jet extending to the northeast in PA=48°. eivine a laree misalignment of APA=112° (?)..," The 5 GHz VLBA map however shows a jet extending to the northeast in $=48\degr$, giving a large misalignment of $\Delta$ $\degr$ \citep{Rector03}."
+ Our VLBI map (Fig. 11)), Our VLBI map (Fig. \ref{fig:1553}) )
+ shows a dominant core and a short jet in PA~45°., shows a dominant core and a short jet in $\sim 45\degr$.
+ There may be fainter jet emission further from, There may be fainter jet emission further from
+chemical enrichment history to each sub-halo that is aecreted by the host galaxy.,chemical enrichment history to each sub-halo that is accreted by the host galaxy.
+ Amone their stated goals is the development of a bimodal color distribution in the GC population without requiring (wo distinct formation mechanisms., Among their stated goals is the development of a bimodal color distribution in the GC population without requiring two distinct formation mechanisms.
+ While the model achieves (his goal by design. it also fits the age-metallicity distribution of the outer halo GCs remarkably well.," While the model achieves this goal by design, it also fits the age-metallicity distribution of the outer halo GCs remarkably well."
+ The lower-right panel of Figure 10. compares the GC AMB with the mean trend (squares) and standard deviation (indicated by error bars) of [Fe/H] for 0.5 Gyr age bins of nearly 10.000 simulated GCs diawn from the Muratov&Gnedin(2010) model.," The lower-right panel of Figure \ref{AMR} compares the GC AMR with the mean trend (squares) and standard deviation (indicated by error bars) of [Fe/H] for 0.5 Gyr age bins of nearly 10,000 simulated GCs drawn from the \citet{mg10} model."
+ While the outer halo GC's are well represented by the mean trend. the model is discrepant with the metal-rich GCs in the inner Galaxy. i.e.. the bulge and thick disk GCs.," While the outer halo GCs are well represented by the mean trend, the model is discrepant with the metal-rich GCs in the inner Galaxy, i.e., the bulge and thick disk GCs."
+ Muratov Gnedin attribute (his discrepancy (o the lack of a metallicity eracient within the individual merging halos., Muratov Gnedin attribute this discrepancy to the lack of a metallicity gradient within the individual merging halos.
+ As a substantial fraction of the Galactic GC population is old aud metal-rich. it behooves a model of the GC population formation to describe the origins of these GCs.," As a substantial fraction of the Galactic GC population is old and metal-rich, it behooves a model of the GC population formation to describe the origins of these GCs."
+ The AMBs constructed by Marin-Franchetal.(2009). ancl Dotteretal.(2010.updatedinthispaper) show that Galactic GCs split into two distinct branches at |Fe/1I] =—1.5. Marfn-, The AMRs constructed by \citet{mf09} and \citet[][further updated in this paper]{do10} show that Galactic GCs split into two distinct branches at [Fe/H] $\ga -1.5$.
+Franchetal.(2009) interpreted this branching as representing (wo groups of GCs: an old group with uniformly old ages (a flat. AMIR) and a voung eroup with a significant (τος (toward vounger ages al higher metalliciüies., \citet{mf09} interpreted this branching as representing two groups of GCs: an old group with uniformly old ages (a flat AMR) and a young group with a significant trend toward younger ages at higher metallicities.
+ Thev concluded (hat these groups likely originated [rom two different phases of Galaxy formation: an imitial. relatively brief collapse and a second. prolonged episode driven by accretion.," They concluded that these groups likely originated from two different phases of Galaxy formation: an initial, relatively brief collapse and a second, prolonged episode driven by accretion."
+ Marin-Eranchetal.(2009) found a small number of metal-poor ([M/II] < —1.5) GCs with somewhat vounger ages (han the bulk of the metal-poor GC's., \citet{mf09} found a small number of metal-poor ([M/H] $\le -1.5$ ) GCs with somewhat younger ages than the bulk of the metal-poor GCs.
+ Dotteretal.(2010) ancl the present study find all of (he metal-poor GCs to be uniformly but this difference does not substantially influence the conclusion with respect to the proposed Galaxy. formation scenario.," \citet{do10}
+ and the present study find all of the metal-poor GCs to be uniformly but this difference does not substantially influence the conclusion with respect to the proposed Galaxy formation scenario."
+ A nmunber of influential studies have used the IB morphologyv-metallicity diagram to inler (he ages of stellar populations for which metallicitv information and good photometry to al least the level of (he IIB are available., A number of influential studies have used the HB morphology-metallicity diagram to infer the ages of stellar populations for which metallicity information and good photometry to at least the level of the HB are available.
+ For example. Catelan&deFreitasPacheco(1993) and Lee.Demarque.&Zinn(1994) both used svnthetic IIB models to demonstrate that. under (he assumption of constant mass loss. the Galactic GC: WB morphology-metallicity diagram showed evidence [or an age spread of several ( 4) Gyr for a fixed He content and heavy element distribution.," For example, \citet{ca93} and \citet{ldz94} both used synthetic HB models to demonstrate that, under the assumption of constant mass loss, the Galactic GC HB morphology-metallicity diagram showed evidence for an age spread of several $\sim4$ ) Gyr for a fixed He content and heavy element distribution."
+ The use of the IIB morphologv-metallicitv diagram must. however. remain a crude," The use of the HB morphology-metallicity diagram must, however, remain a crude"
+observations was estimated to be 16.541.6 Jy (Richards. private communication).,"observations was estimated to be $\pm$ 1.6 Jy (Richards, private communication)."
+ The phase calibrator source was mapped with one round of phase self-calibration followed by an amplitude and phase self calibration., The phase calibrator source was mapped with one round of phase self-calibration followed by an amplitude and phase self calibration.
+ Again IMAGR was used to make the Images and CLEAN them., Again IMAGR was used to make the images and CLEAN them.
+ The resulting FWHM of the beam is 26 x 24 mas at a position angle of 397., The resulting FWHM of the beam is 26 $\times$ 24 mas at a position angle of $39^\circ$.
+ The accuracy of the absolute masers position measured in the paper is limited by four factors: (1) the position accuracy of the phase calibrator. (2) the accuracy of the telescope positions. (3) the relative position error depending on the beamsize and signal-to-noise ratio and (4) finally the atmospheric variability that plays an important role especially at 22 GHz depending on the angular separation between the calibrator (20054403) and the target.," The accuracy of the absolute masers position measured in the paper is limited by four factors: (1) the position accuracy of the phase calibrator, (2) the accuracy of the telescope positions, (3) the relative position error depending on the beamsize and signal-to-noise ratio and (4) finally the atmospheric variability that plays an important role especially at 22 GHz depending on the angular separation between the calibrator (2005+403) and the target."
+ The first two factors are frequency independent and were estimated to be 5 mas (given by the MERLIN calibrator catalogue) and 10 mas respectively (Diamond et al., The first two factors are frequency independent and were estimated to be 5 mas (given by the MERLIN calibrator catalogue) and 10 mas respectively (Diamond et al.
+ 2003)., 2003).
+ The other two factors are frequency dependent., The other two factors are frequency dependent.
+ The relative position error. given approximately by the beamsize/signal-to-noise ratio. leads to uncertainties in the position of 17. 4 and 2.5 mas at OH. aand mmaser lines respectively.," The relative position error, given approximately by the beamsize/signal-to-noise ratio, leads to uncertainties in the position of 17, 4 and 2.5 mas at OH, and maser lines respectively."
+ Factor (4) is inferred from. the quality of the phase of the calibrator., Factor (4) is inferred from the quality of the phase of the calibrator.
+ With a separation of 2° between the phase calibrator and the target this factor adds an uncertainties in the absolute position of 5 to 20 mas at the 1665-MHz OH maser line for the typical and worst phase rate respectively., With a separation of $2^\circ$ between the phase calibrator and the target this factor adds an uncertainties in the absolute position of 5 to 20 mas at the 1665-MHz OH maser line for the typical and worst phase rate respectively.
+ For the aand 6.7-GHz mmaser lines. it adds errors of 10 and 2 mas respectively. taking into account the worst phase rate for each observation.," For the and 6.7-GHz maser lines, it adds errors of 10 and 2 mas respectively, taking into account the worst phase rate for each observation."
+ All these uncertainties combine quadratically to give absolute position errors of 25. 15 and 12 mas in the 1665-MHz OH. 22-GHz aand 6.7-GHz ccomponent positions respectively.," All these uncertainties combine quadratically to give absolute position errors of 25, 15 and 12 mas in the 1665-MHz OH, 22-GHz and 6.7-GHz component positions respectively."
+ The 1665-MHz OH line. water and 6.7-GHz methanol masers were detected with. MERLIN.," The 1665-MHz OH line, water and 6.7-GHz methanol masers were detected with MERLIN."
+ Table 2 gives the absolute positions and. velocities of the brightest maser spot for each maser type., Table 2 gives the absolute positions and velocities of the brightest maser spot for each maser type.
+ Radial velocities. here and elsewhere. are given relative to the Local Standard of Rest (LSR).," Radial velocities, here and elsewhere, are given relative to the Local Standard of Rest (LSR)."
+ All three types of maser are located close to the central source., All three types of maser are located close to the central source.
+ The details of the emission in each maser type are discussed below., The details of the emission in each maser type are discussed below.
+ The MHz OH maser line was not detected down to a noise level of 20 mJy/beam., The 1667-MHz OH maser line was not detected down to a noise level of 20 mJy/beam.
+ A total of 9 (7 RHC and 2 LHC) 1665-MHz OH maser spots were detected towards IRAS 20126+4104., A total of 9 (7 RHC and 2 LHC) 1665-MHz OH maser spots were detected towards IRAS $20126+4104$.
+ Table 3 presents the parameters of the OH maser components detected. namely the peak intensities. velocities and positions for each hand of circular polarisation.," Table 3 presents the parameters of the OH maser components detected, namely the peak intensities, velocities and positions for each hand of circular polarisation."
+ The label Z marks a left-hand and right-hand polarised pair of components which originate at the same location and are identified as a Zeeman pair with a splitting of 6.3 km s!., The label Z marks a left-hand and right-hand polarised pair of components which originate at the same location and are identified as a Zeeman pair with a splitting of 6.3 km $^{-1}$.
+ Figure 1. shows the distribution of the OH maser spots., Figure \ref{fig:oh_radio_position} shows the distribution of the OH maser spots.
+ They are distributed 1n two clusters on opposite sides of the axis inferred for the radio continuum jet (C99. Hofner et al.," They are distributed in two clusters on opposite sides of the axis inferred for the radio continuum jet (C99, Hofner et al."
+ 1999) and about 0.3” south-east from the 3.6em continuum centre., 1999) and about $0.3''$ south-east from the 3.6cm continuum centre.
+ The OH maser spots are spread over a region of ~1.2” corresponding to 2000 AU (or ~ 0.01 pe) at a distance 1.7 kpe., The OH maser spots are spread over a region of $\sim1.2''$ corresponding to 2000 AU (or $\sim$ 0.01 pc) at a distance 1.7 kpc.
+ The distribution of OH masers is approximately symmetrical about a line of NE-SW direction at position angle 29°., The distribution of OH masers is approximately symmetrical about a line of NE-SW direction at position angle $^\circ$.
+ Figure | also shows the spatial distribution of the velocities of the maser components., Figure \ref{fig:oh_radio_position} also shows the spatial distribution of the velocities of the maser components.
+ For the OH masers the velocities are mainly negative to the north and positive to the south., For the OH masers the velocities are mainly negative to the north and positive to the south.
+ Table 4 present the Stokes parameters 1. Q. U and V. the polarisation position angle (y) (angles are measured from towards E). the linearly. polarised flux P. the percentage of linear polarisation 775. the percentage of circular polarisatior mc and the total percentage of polarisation zy of each feature.," Table 4 present the Stokes parameters I, Q, U and V, the polarisation position angle $\chi$ ) (angles are measured from N towards E), the linearly polarised flux P, the percentage of linear polarisation $m_L$, the percentage of circular polarisation $m_C$ and the total percentage of polarisation $m_T$ of each feature."
+ The OH masers have a total polarisation around ranging from to100., The OH masers have a total polarisation around ranging from to.
+.. The Stokes intensities are shown as zero i1 this table if their flux is below the noise level., The Stokes intensities are shown as zero in this table if their flux is below the noise level.
+ All 1665-MHz features are circularly polarised and three features (3. 5 anc 6) are elliptically polarised.," All 1665-MHz features are circularly polarised and three features (3, 5 and 6) are elliptically polarised."
+ Feature 3 is the most elliptically polarised (36.4%)) and feature 7 1s 100 circularly polarised., Feature 3 is the most elliptically polarised ) and feature 7 is 100 circularly polarised.
+ Polarisation position angles can only be measured for features 3 and 5., Polarisation position angles can only be measured for features 3 and 5.
+" The water masers towards IRAS 20126+4104 were detectec in a single cluster of ~0.15"" (255 AU at 1.7 kpe) in size close to the central source and the OH (and methanol. 3.39) masers."," The water masers towards IRAS $20126+4104$ were detected in a single cluster of $\sim0.15''$ (255 AU at 1.7 kpc) in size close to the central source and the OH (and methanol, \ref{sec:meth}) ) masers."
+ This cluster is close to the C2 group of spots identified by Tofani et al. (, This cluster is close to the C2 group of spots identified by Tofani et al. (
+1995).,1995).
+ However the positions and fluxes have varied strongly. the brightest emission being 20 times greater than seen by Tofani et al.," However the positions and fluxes have varied strongly, the brightest emission being 20 times greater than seen by Tofani et al."
+ Also. neither of the clusters CI and C3 identified by Tofani et al.," Also, neither of the clusters C1 and C3 identified by Tofani et al."
+ was detected with MERLIN. implying that they have decreased in brightness by factors of more than 10 and 100 respectively.," was detected with MERLIN, implying that they have decreased in brightness by factors of more than 10 and 100 respectively."
+ Figure 2 shows that the spectrum of the water masers detected with MERLIN ts also considerably different from that observed by MCR., Figure \ref{fig:22spec} shows that the spectrum of the water masers detected with MERLIN is also considerably different from that observed by MCR.
+ The strongest component detected by MCR centred at Vise ~12 km s! is not detected with MERLIN nor are the highest velocity components centred between -30 km s! and 20 km s'., The strongest component detected by MCR centred at $V_{LSR}$ $\sim 12$ km $^{-1}$ is not detected with MERLIN nor are the highest velocity components centred between -30 km $^{-1}$ and 20 km $^{-1}$.
+ On the other hand. MERLIN detected components between 0 km s! and +5 km s7! which were previously not detected with the VLBA.," On the other hand, MERLIN detected components between 0 km $^{-1}$ and +5 km $^{-1}$ which were previously not detected with the VLBA."
+ At the same time. the component centred in the velocity range -5.0 km s! aand 0 km s! has increased in strength and the component centred at ~ -7kms7! dominates both spectra. although as Figure 3 shows. the location of the emission at this velocity is significantly different between the two observations.," At the same time, the component centred in the velocity range -5.0 km $^{-1}$ and 0 km $^{-1}$ has increased in strength and the component centred at $\sim-7$ km $^{-1}$ dominates both spectra, although as Figure \ref{fig:h2o_comp} shows, the location of the emission at this velocity is significantly different between the two observations."
+"As has been shown by Braun et ((2009), M31 has an disk that extends to radii of r~30 kpc for column densities above log 20.3.","As has been shown by Braun et (2009), M31 has an disk that extends to radii of $r\approx 30$ kpc for column densities above log $N$ $)\geq 20.3$."
+ Beyond r=30 kpc the disk appears to be N(H1)>truncated (Braun et 22009; Braun Thilker 2004)., Beyond $r\approx 30$ kpc the disk appears to be truncated (Braun et 2009; Braun Thilker 2004).
+ The HVC population of M31 was studied in detail with the Green Bank Telescope (GBT) by Thilker et ((2004)., The HVC population of M31 was studied in detail with the Green Bank Telescope (GBT) by Thilker et (2004).
+" These authors detected more than 20 individual HVCs around M31 with column densities >5x1017 cm-? and estimated a total mass of the HVCs of 3—4x10""Mo (see also Thilker, Braun Westmeier 2005)."," These authors detected more than 20 individual HVCs around M31 with column densities $\geq 5\times10^{17}$ $^{-2}$ and estimated a total mass of the HVCs of $3-4\times 10^7\,M_{\sun}$ (see also Thilker, Braun Westmeier 2005)."
+" In the left panel of 11 we show the distribution of HVCs around the M31 disk, based on the GBT 21cm contour map presented by Thilker et "," In the left panel of 1 we show the distribution of HVCs around the M31 disk, based on the GBT 21cm contour map presented by Thilker et (2004)."
+The HVC population of M31 reaches out until ~50 kpc al.((2004).with a strongly decreasing HVC covering fraction towards larger radii., The HVC population of M31 reaches out until $\sim 50$ kpc with a strongly decreasing HVC covering fraction towards larger radii.
+" We assume D=785+25 kpc as distance for M31 (McConnachie et al.22005), so that absolute distance estimates derived from angular coordinates are uncertain by ~3 percent."," We assume $D=785\pm25$ kpc as distance for M31 (McConnachie et 2005), so that absolute distance estimates derived from angular coordinates are uncertain by $\sim 3$ percent."
+ Some interesting conclusions about the radial distribution of neutral gas around M31 can be drawn from this HVC distribution map., Some interesting conclusions about the radial distribution of neutral gas around M31 can be drawn from this HVC distribution map.
+" 11, right panel, shows the radius-dependent projected covering fraction, of the HVC population of M31 plotted against r fuvc(r), by the filled where r is the projected (indicatedradius."," 1, right panel, shows the radius-dependent projected covering fraction, $f_{\rm HVC}(r)$, of the HVC population of M31 plotted against $r$ (indicated by the filled boxes), where $r$ is the projected radius."
+" To calculate boxes),fuvc(r) we have resampled the M31 HVC map of Thilker et ((2004) and have transformed the coordinate system into a polar coordinate system (o,with0) coordinates r and 0 centered on M31."," To calculate $f_{\rm HVC}(r)$ we have resampled the M31 HVC map of Thilker et (2004) and have transformed the $\alpha,\delta$ ) coordinate system into a polar coordinate system with coordinates $r$ and $\theta$ centered on M31."
+ For each ring with radius r and thickness r+Ar the parameter fuvc(r) then was derived by comparing the area covered by HVC gaswith the total ring area., For each ring with radius $r$ and thickness $r\pm\Delta r$ the parameter $f_{\rm HVC}(r)$ then was derived by comparing the area covered by HVC gaswith the total ring area.
+ The error bars for fuvyc(r) shown in 11 have been calculated assuming Poisson-like statistics., The error bars for $f_{\rm HVC}(r)$ shown in 1 have been calculated assuming Poisson-like statistics.
+ Starting from fuvc70.5 at r=15 kpc the covering fraction decreases to values less than 0.05 for radii larger than r=45 kpc., Starting from $f_{\rm HVC}\approx 0.5$ at $r=15$ kpc the covering fraction decreases to values less than $0.05$ for radii larger than $r=45$ kpc.
+" This trend for favc(r) can be fitted by an exponential in the form fuvc(r)=xeexp (—r/y) with x=2.10.2 and y= as shown by the solid red line in the right panel12.0492, of 11."," This trend for $f_{\rm HVC}(r)$ can be fitted by an exponential in the form $f_{\rm HVC}(r)=x$ $\,(-r/y)$ with $x=2.1\pm0.2$ and $y=12.0^{+0.7}_{-0.5}$, as shown by the solid red line in the right panel of 1."
+" An exponential fit to the extraplanar features of M31 was also favored by Braun Thilker (2004), who analyzed lower-resolution 21cm data of M31 from theTelescope (WSRT)."," An exponential fit to the extraplanar features of M31 was also favored by Braun Thilker (2004), who analyzed lower-resolution 21cm data of M31 from the (WSRT)."
+ For the inner regions of the M31 halo at r«15 kpc no HVC data are available., For the inner regions of the M31 halo at $r<15$ kpc no HVC data are available.
+" As discussed by Thilker et al.((2004), this does not imply that this region is devoid of HVC material."," As discussed by Thilker et (2004), this does not imply that this region is devoid of HVC material."
+" The lack of data for the disk-halo interface region at r«15 kpc left panel) rather indicates the incompleteness (Fig.11,of the HVC map for small radii because of the confusion of neutral halo gas with the disk of M31 together with the stringent selection criteria defined by Thilker et to unambiguously identify HVC features."," The lack of data for the disk-halo interface region at $r<15$ kpc 1, left panel) rather indicates the incompleteness of the HVC map for small radii because of the confusion of neutral halo gas with the disk of M31 together with the stringent selection criteria defined by Thilker et to unambiguously identify HVC features."
+" In the Milky Way, the disk-halo interface at r«15 kpc is filled with large amounts of neutral gas that gives rise to 21cm emission at intermediate and high velocities (e.g., Wakker 2004)."," In the Milky Way, the disk-halo interface at $r<15$ kpc is filled with large amounts of neutral gas that gives rise to 21cm emission at intermediate and high velocities (e.g., Wakker 2004)."
+ Nearby edge-on galaxies such as 8891 also exhibit large amounts of neutral gas in the disk-halo interface region extending several kpc above and below the disk (see Sancisi et 2, Nearby edge-on galaxies such as 891 also exhibit large amounts of neutral gas in the disk-halo interface region extending several kpc above and below the disk (see Sancisi et 2008).
+1cm measurements of NGC8891 show that the al.22008).(projected) covering fraction of neutral gas is fui=1 for vertical distances d<10 kpc to the midplane of the 8891 disk (Oosterloo et 22007)., 21cm measurements of 891 show that the (projected) covering fraction of neutral gas is $f_{\rm HI}=1$ for vertical distances $d<10$ kpc to the midplane of the 891 disk (Oosterloo et 2007).
+" If we extrapolate fuvc for M31 to small radii using the exponential defined above, /ηνο becomes unity for r<9 kpc, in line with the extraplanar gas distribution observed in 8891."," If we extrapolate $f_{\rm HVC}$ for M31 to small radii using the exponential defined above, $f_{\rm HVC}$ becomes unity for $r\leq 9$ kpc, in line with the extraplanar gas distribution observed in 891."
+" If we define r3 as the radius beyond which the projected covering fraction of HVCs falls below 3 percent, we derive for M31 a value of T3=50+6 kpc."," If we define $r_3$ as the radius beyond which the projected covering fraction of HVCs falls below 3 percent, we derive for M31 a value of $r_3=50\pm 6$ kpc."
+ The 1o error reflects the uncertainties in the exponential parameters for fuvc(r)., The $1\sigma$ error reflects the uncertainties in the exponential parameters for $f_{\rm HVC}(r)$ .
+" As shown by Braun Thilker (2004), the high-resolution 21cm HVC data of M31 from Thilker et al.((2004) follows a “standard”"," As shown by Braun Thilker (2004), the high-resolution 21cm HVC data of M31 from Thilker et (2004) follows a “standard”"
+account the error on Ze.,account the error on $R_V$.
+ We were to caleulate the host galaxy extinction for GRB050822 nodsince there is no information about its host -extinction. profile in the literature., We were unable to calculate the host galaxy extinction for GRB 050822 since there is no information about its host galaxy extinction profile in the literature.
+ We did not calculate Galactic extinction for GIU 050822. GIUD 060729. GARB 060927. and CRB 061222 since all of the included optical [Dux densities for these bursts have already been corrected for Galactic extinction by either Yost et al. (," We did not calculate Galactic extinction for GRB 050822, GRB 060729, GRB 060927, and GRB 061222A since all of the included optical flux densities for these bursts have already been corrected for Galactic extinction by either Yost et al. ("
+200Ta.b) or Ruiz-Velasco et al. (,"2007a,b) or Ruiz-Velasco et al. ("
+2007).,2007).
+ Phe errors on ely. were negligible., The errors on $A_{\lambda_{obs}}$ were negligible.
+ The ΝΤΕ data for the in our sample are taken from the data, The /XRT data for the GRBs in our sample are taken from the data.
+ standardCRBsWe developed an IDL package which invokes the οΡΕ tools (c.g.. Nselect. Nimage. Nspec. Swift tools) to automatically oocess the NAP data for anv eiven burst.," We developed an IDL package which invokes the standard HEASoft tools (e.g., Xselect, Ximage, Xspec, Swift tools) to automatically process the XRT data for any given burst."
+ The details o£ he procedure of our package are described in Zhang ct al. , The details of the procedure of our package are described in Zhang et al. (
+2007b).,2007b).
+ Our final data include the NIE. light curves (in physical units). extracted over the energy. range L3.LOkeV.," Our final data include the XRT light curves (in physical units), extracted over the energy range $0.3-10~keV$."
+ These cata products have been applied. to study statistical properties of X-ray afterglows in a series of xipers (Zhang et al., These data products have been applied to study statistical properties of X-ray afterglows in a series of papers (Zhang et al.
+ 2007b: Liang et al., 2007b; Liang et al.
+ 2007. 2008. 2009).," 2007, 2008, 2009)."
+ For the bursts in our sample. the spectral. index j for the ART band is obtained by Ditting the time-integrated spectrum. (Evans et al.," For the bursts in our sample, the spectral index $\beta$ for the XRT band is obtained by fitting the time-integrated spectrum (Evans et al."
+ 2009) of the shallow and norma decay phases with a simple »ower Law mociel in Xspec., 2009) of the shallow and normal decay phases with a simple power law model in Xspec.
+" The spectral fit also. includes: two-component neutral hydrogen column density Vyíi, and μοι with corresponding values taken from the late-time spectral data in the Swift/XIT GRB spectrum (Evans οἱ al."," The spectral fit also includes two-component neutral hydrogen column density $N_{H,host}$ and $N_{H,Gal}$, with corresponding values taken from the late-time spectral data in the Swift/XRT GRB spectrum (Evans et al."
+ 2009)., 2009).
+ Their values were held constant in our spectral fit. which are collected in Table 1 along with the GRB redshift > and the spectral fitting results.," Their values were held constant in our spectral fit, which are collected in Table 1 along with the GRB redshift $z$ and the spectral fitting results."
+ All the errors quoted in this paper correspond to the le errors on the quantities. unless otherwise indicated.," All the errors quoted in this paper correspond to the $1\sigma$ errors on the quantities, unless otherwise indicated."
+ The convention f.[o5 Cds adopted., The convention $F_{\nu}\varpropto{ }t^{-\alpha}\nu^{-\beta}$ is adopted.
+ The zero points aud elfective frequencies for the optical bands from Dessell ( are adopted., The zero points and effective frequencies for the optical bands from Bessell (1979) are adopted.
+ In this paper we focus on the NIE data in the shallow and normal decay phases (clark black points in Fig., In this paper we focus on the XRT data in the shallow and normal decay phases (dark black points in Fig.
+ 3)., 3).
+ We do not consider the post jet break phase as the jet break happens much later., We do not consider the post jet break phase as the jet break happens much later.
+ Using the prior emission model and Eq. 2]], Using the prior emission model and Eq. \ref{X2}] ]
+ from Liang ct al. (, from Liang et al. (
+2t09) as our prescription. we the shallow and normal decay segments of cach NICE light curve to the function allowing the normalization constant fy. time shift constant Z4. and temporal decay index a to vary in the fit.,"2009) as our prescription, we the shallow and normal decay segments of each XRT light curve to the function allowing the normalization constant $F_0$, time shift constant $T_\Delta$, and temporal decay index $\alpha$ to vary in the fit."
+ Here T£ is thie ‘time measured since the GRB trigger., Here $T$ is the time measured since the GRB trigger.
+ The fitting intervials are loosely based on the time ranges eiven in Table 1 of Liang et al. (, The fitting intervals are loosely based on the time ranges given in Table 1 of Liang et al. (
+2007).,2007).
+ Once the parameters have been determined for à particular burst. Eq. 1]]," Once the parameters have been determined for a particular burst, Eq. \ref{X1}] ]"
+ can be used to extrapolae the X-ray flux due to the prior emission component [rom the shallow decay segment to earlier times (Le. to the time of the early optical observations of interest).," can be used to extrapolate the X-ray flux due to the prior emission component from the shallow decay segment to earlier times (i.e., to the time of the early optical observations of interest)."
+ The results of the XR light curve fits are given in Table 2., The results of the XRT light curve fits are given in Table 2.
+ The external shock model predicts a broken power law synchrotron προςrum (Alésszarros Rees 1997 Sark οἱ al., The external shock model predicts a broken power law synchrotron spectrum (Mésszárros Rees 1997; Sari et al.
+ 1998)., 1998).
+" For à constant clensity medium. a fast cooling spectrum (4.« £u) is expected. at carly times. which transitions to asow cooling spectrum (η, £a) at later times (Sari ct al."," For a constant density medium, a fast cooling spectrum $\nu_c < \nu_m$ ) is expected at early times, which transitions to a slow cooling spectrum $\nu_m < \nu_c$ ) at later times (Sari et al."
+ 1998)., 1998).
+ In our analysis. since Z4~10105s (tvpical obseved. shallow-to-normal break time). the prior emission exernal shock should have already entered the slow cooling regime by the time of the GRB trigger.," In our analysis, since $T_\Delta\sim10^3-10^4~s$ (typical observed shallow-to-normal break time), the prior emission external shock should have already entered the slow cooling regime by the time of the GRB trigger."
+ Thus. we apply slow cooling throughout the modeling.," Thus, we apply slow cooling throughout the modeling."
+ supernovae have emerged as the most promising standard candles.,} Supernovae have emerged as the most promising standard candles.
+ Due (o their significant intrinsic brightness and relative ubiquity (μεν can be observed in the local and distant universe., Due to their significant intrinsic brightness and relative ubiquity they can be observed in the local and distant universe.
+ Observational efforts to detect high-redshift supernovae have proved their value as cosmological probes., Observational efforts to detect high-redshift supernovae have proved their value as cosmological probes.
+ The systematic study ancl observation of these faint supernovae (mainly ivpe Ia) has been utilized to constrain the cosmic expansion history (Goobar&PerlinutterPerlnutteretal.1999:Schmidtοἱ 1993).," The systematic study and observation of these faint supernovae (mainly type Ia) has been utilized to constrain the cosmic expansion history \citep{Goobar95, Perlmutter99, Schmidt98}."
+. Light emitted from anv celestial object is subject. to lensing bv intervening objects while traversing the large distances involved (Ixantowski.Vaughan.&Braneh1995) and the farther the light source. the higher its chance of being signilicantly lensed.," Light emitted from any celestial object is subject to lensing by intervening objects while traversing the large distances involved \citep{KVB95} and the farther the light source, the higher its chance of being significantly lensed."
+ Apart from the fact that eravitational lensing can limit the accuracy of luminosity clistance measurements (Perlmutter&Schmidi2003).. it can change the observed rate of supernovae as well.," Apart from the fact that gravitational lensing can limit the accuracy of luminosity distance measurements \citep{Perlmutter03}, it can change the observed rate of supernovae as well."
+ Slucdving supernovae and their rates at high redshifis provide us with much needed information for constraining the measurements of the ellusive dark energy. as well as understanding the cosmic star formation rate and metal enrichment at high vedshifts.," Studying supernovae and their rates at high redshifts provide us with much needed information for constraining the measurements of the ellusive dark energy, as well as understanding the cosmic star formation rate and metal enrichment at high redshifts."
+ In order to observe and. hence. study (he faint high-redshift supernovae. one can raise (he chance of observation by looking through clusters of galaxies or even massive galaxies (see 9mail.etal.(2002) and (he references therein).," In order to observe and, hence, study the faint high-redshift supernovae, one can raise the chance of observation by looking through clusters of galaxies or even massive galaxies (see \citet{Sma02} and the references therein)."
+ These ‘eravilational telescopes’ amplify the high-redshift supernovae and thereby increase the chance of their detection., These `gravitational telescopes' amplify the high-redshift supernovae and thereby increase the chance of their detection.
+ However. this boost in observation is olfset bv the competing effect of depletion (Fig.," However, this boost in observation is offset by the competing effect of depletion (Fig."
+ 1). due to the Ποιά being spread by the deflector (amplification bias).," 1), due to the field being spread by the deflector (amplification bias)."
+ For an assumed lens model ancl a given field of view it is nol obvious which effect. dominates the observation of supernovae (through (he halo., For an assumed lens model and a given field of view it is not obvious which effect dominates the observation of supernovae through the halo.
+ The net result depends on the deflector and source parametrs as well as the observational setup (Gunnarsson&Goobar2003)., The net result depends on the deflector and source parametrs as well as the observational setup \citep{GunnGoo03}.
+. some research has been conducted on the feasibility of observing supernovae through cluster of galaxies, Some research has been conducted on the feasibility of observing supernovae through cluster of galaxies
+Observatiois Πρίν that there maybe a superiassive black hole at the ceuter of alinost every galaxies.,Observations imply that there maybe a supermassive black hole at the center of almost every galaxies.
+ There ix a large raction for which the accretion flow can be described. by a hot accretion flow. such as advectiou-dominated accretion flow(ADAF: Naravan Yi 1991. 1995: Abranowicz et al.," There is a large fraction for which the accretion flow can be described by a hot accretion flow, such as advection-dominated accretion flow (ADAF; Narayan Yi 1994, 1995; Abramowicz et al."
+ 1995: for a review of ADAE. see Narayan. \Tahadevan Quataert 1998: for a review of he applicaious of ADAF. see Yuan 2007) ancl luminous rot aceretion flow (LITAF: Yuan 2001).," 1995; for a review of ADAF, see Narayan, Mahadevan Quataert 1998; for a review of the applications of ADAF, see Yuan 2007) and luminous hot accretion flow (LHAF; Yuan 2001)."
+ Tot accretion flow also ¢ xdbstsd1 sonue states of black hole N-ray |iuaries., Hot accretion flow also exists in some states of black hole X-ray binaries.
+ The claperature of the hot Bow is- very high.B 7;~ELot?D ir. F clue the racius in uuit of Sclawarzsclild radius ry.," The temperature of the hot flow is very high, $
+T_{i}\sim10^{12}$ $/r$, $r$ being the radius in unit of Schwarzschild radius $r_s$."
+ This cliperature is high chough to ignite sigUficant uuclear reactions iu the hot flow., This temperature is high enough to ignite significant nuclear reactions in the hot flow.
+ This process is poteutia iuteresting because it is a pronuising nieaas of produci jew. elements., This process is potentially interesting because it is a promising means of producing new elements.
+ The aalysis of observatious by JiALCZ al. (, The analysis of observations by Juarez et al. (
+2109: sce also Tatain Forlanc 1999) found that netalicity of the DER o ‘redshift quasars is very higLi.,2009; see also Hamann Ferland 1999) found that the metallicity of the BLR of redshift quasars is very high.
+ ο) explanation is that t10 lnassive stars had formic before he quasars were created. and these sars then euricied 1ο netalicities of quasars.," One explanation is that the massive stars had formed before the quasars were created, and these stars then enriched the metallicities of quasars."
+ It is also possiblethat DLR eas origilates In accretion flows., It is also possible that BLR gas originates in accretion flows.
+ Although accretion flowsin quasars are eenerallv assumed to beii the regine o anmlor accrelon rates than that of a hot accretiou flow. Le. he Stauard thin disk. the detailed structure of he accretion flow is still largely uuknown because the thin disk can iot explain the origin of the N-rav. eiiissiou widely observed iu qiasars.," Although accretion flows in quasars are generally assumed to be in the regime of higher accretion rates than that of a hot accretion flow, i.e., the standard thin disk, the detailed structure of the accretion flow is still largely unknown because the thin disk can not explain the origin of the X-ray emission widely observed in quasars."
+" The presence of X-ray enission muüplies tlat ADAF or LITAF-like hot accretion flow is also Likely to exist I quasars. uot only in some low-Inuinositv Αννα,"," The presence of X-ray emission implies that ADAF or LHAF-like hot accretion flow is also likely to exist in quasars, not only in some low-luminosity AGNs."
+ One oossible reason for the high metallicity in quasars is therefore that uucleosvuthesis occurs in the accretion flow (see the review paper of Ibuuaun Ferland 1999)., One possible reason for the high metallicity in quasars is therefore that nucleosynthesis occurs in the accretion flow (see the review paper of Hamann Ferland 1999).
+" Some people have already studied the nuclear reactions aud 1imcleosvuthesis in accretion flows around]ack. holes (ο,ο, Jin et al."," Some people have already studied the nuclear reactions and nucleosynthesis in accretion flows around black holes (e.g., Jin et al."
+ 1989. Arii IIshimoto 1992. Chakrabarti AIuk10]hvay 1999. Mikhopadhivay Chakrabarti 2000. Mikhopadhivay Chakrabarti 2001. Tn Peue 2008).," 1989, Arai Hashimoto 1992, Chakrabarti Mukhopadhyay 1999, Mukhopadhyay Chakrabarti 2000, Mukhopadhyay Chakrabarti 2001, Hu Peng 2008)."
+ Iu these works. however. the accretion flow models adopted are cite different from ADAFs.," In these works, however, the accretion flow models adopted are quite different from ADAFs."
+ The works of Jiu et al., The works of Jin et al.
+ and Arai Tashinoto are based on a “thick disk” while the two works by Mukhliopadlivay Chakrabarti are based ou a siuplifted inviscid accreion flow. which differs aeain frou ADAFs.," and Arai Hashimoto are based on a “thick disk”, while the two works by Mukhopadhyay Chakrabarti are based on a simplified inviscid accretion flow, which differs again from ADAFs."
+ These particular accretion models have receive’ little atteition these vears., These particular accretion models have received little attention these years.
+ Tu Peng (2008) studied uucleosvuthesis in ADAFs. but their work was based on the selfsimular solution of ADAF.," Hu Peng (2008) studied nucleosynthesis in ADAFs, but their work was based on the self-similar solution of ADAF."
+ This solutilon is affected by significaut error iu the ier region of the accretion flow. where most nuclear τςactions occur.," This solution is affected by significant error in the inner region of the accretion flow, where most nuclear reactions occur."
+ We herefore revisit the «παν of nuccosvuthesis iu hot accrejon flows., We therefore revisit the study of nucleosynthesis in hot accretion flows.
+ We first coustrain the νταides of accretion flows., We first constrain the dynamics of accretion flows.
+ Althougi the global solution of ADAE has heen |-iowii for nuiuiv vears. since nuclear reactions have nuo 'eviouslv been taken iuto accouit. we reed to Investigate whether they are important compared tot 1C viscous lie:ine.," Although the global solution of ADAF has been known for many years, since nuclear reactions have not previously been taken into account, we need to investigate whether they are important compared to the viscous heating."
+" If they are then. the seltf-cousisteu elolval solution o"" ADAF needs to be recalculated."," If they are then, the self-consistent global solution of ADAF needs to be recalculated."
+ We negect the photodisiuteeration because the optical deoth in ADAFs is very small aud most of the photos cau escape froni the hot accretion flow without οσο scattered., We neglect the photodisintegration because the optical depth in ADAFs is very small and most of the photos can escape from the hot accretion flow without being scattered.
+ The sructure of the paper is as follows., The structure of the paper is as follows.
+ We Irietiv introduce ADAFs in Sect., We briefly introduce ADAFs in Sect.
+ 2.1 and the caleulation method of energv generation rate in Sect., 2.1 and the calculation method of energy generation rate in Sect.
+ 2.2., 2.2.
+ Ihe calculation results will be shown in Sect., The calculation results will be shown in Sect.
+ 23., 2.3.
+ InSoct., In Sect.
+ ) woe sunnarize our results., 3 we summarize our results.
+ The οναος of the ADAF js described 1»* the continuity equation. the radial aud azimuthal com20Ο15 of momenta equation. and the energev equations of electrons and ious (Naravan. Maliacdevau," The dynamics of the ADAF is described by the continuity equation, the radial and azimuthal components of momentum equation, and the energy equations of electrons and ions (Narayan, Mahadevan"
+enission Components visible in the pulsating tail (see Fig. 3)),emission components visible in the pulsating tail (see Fig. \ref{fol757-aver}) )
+ suggests separate emission regions., suggests separate emission regions.
+ Ht is interesting to note that the pulsations are present since (he very early phase of the flare. as indicated by the broad pulse visible αἱ (=25 s (Fie. 2)).," It is interesting to note that the pulsations are present since the very early phase of the flare, as indicated by the broad pulse visible at t=2–5 s (Fig. \ref{initialfit}) )."
+ This pulse persists al (he same rotational phase in the following cycles. while a secondary. peak al o~0.75 gradually emerges.," This pulse persists at the same rotational phase in the following cycles, while a secondary peak at $\phi\sim$ 0.75 gradually emerges."
+ While the above results are similar to those obtained with other satellites. our data provide unique evidence for a second. long lasting component in the hard Xraw light curve.," While the above results are similar to those obtained with other satellites, our data provide unique evidence for a second, long lasting component in the hard X–ray light curve."
+ The emission alter (+400 s could originate from the neutron star surface and/or maegnetosphere or. alternativelv. [rom (he interaction of the relativistc fireball with the surrounding medium.," The emission after $\sim$ 400 s could originate from the neutron star surface and/or magnetosphere or, alternatively, from the interaction of the relativistic fireball with the surrounding medium."
+ A search for pulsations in the ACS data after t—400 s gave a negative result. lavoring the second interpretation.," A search for pulsations in the ACS data after t=400 s gave a negative result, favoring the second interpretation."
+ The alterelow produced by the egiant [lare has been observed in the radio band (Gaensler et al., The afterglow produced by the giant flare has been observed in the radio band (Gaensler et al.
+ 2005: Cameron οἱ al., 2005; Cameron et al.
+ 2005). vielding estimates of the minimum energy in magnetic field and relativistic particles. EZ. of several LOM erg (see also Nakar. Piran Sari 2005).," 2005), yielding estimates of the minimum energy in magnetic field and relativistic particles, E, of several $^{43}$ erg (see also Nakar, Piran Sari 2005)."
+ The hard. X.rav (hence in the 400-3000 s time interval is consistent. with (his value., The hard X–ray fluence in the 400-3000 s time interval is consistent with this value.
+ With simple gamma-ray burst afterglow models based on svuchrotvon emission we can estimate the bulk Lorentz [actor D. [rom the time t4 of the afterglow onset: D—15 [E / 5 107 ere)? (n / 0.1 em?) |? (t / 1008) 7. where n is the ambient density. (see. e.g. Zhang Meszaros 2003. and references (herein).," With simple gamma-ray burst afterglow models based on synchrotron emission we can estimate the bulk Lorentz factor $\Gamma$ from the time $_{0}$ of the afterglow onset: $\Gamma\sim$ 15 (E / 5 $^{43}$ $^{1/8}$ (n / 0.1 $^{-3}$ $^{-1/8}$ $_{0}$ / 100 $^{-3/8}$, where n is the ambient density (see, e.g. Zhang Meszaros 2003, and references therein)."
+ This is smaller (han the tvpical values or gamma-ray bursts. but consistent. considering the large involved uncertainties. with the wildly relativistic outflow inferred from the modelling of the radio data (Granot et al.," This is smaller than the typical values for gamma-ray bursts, but consistent, considering the large involved uncertainties, with the mildly relativistic outflow inferred from the modelling of the radio data (Granot et al."
+ 2005)., 2005).
+ Alternatively. the observed re-brightening could be due to an Inverse Compton component. implying a hieh density environment 2 10 o a and a high electron radiation efficiency.," Alternatively, the observed re-brightening could be due to an Inverse Compton component, implying a high density environment $\gsim$ 10 $^{-3}$ ) and a high electron radiation efficiency."
+ We thank Ix. IIurley. G. Ghisellini. J. Borkowski and M. Feroci for useful comments.," We thank K. Hurley, G. Ghisellini, J. Borkowski and M. Feroci for useful comments."
+ This work has been partially funded by the Italian Space Agency., This work has been partially funded by the Italian Space Agency.
+" The SPI-ACS is supported by ihe German Ministerium fur Bildung und Forschung"" through the DLR grant 50.0G.9503.0."," The SPI-ACS is supported by the German ""Ministerium fur Bildung und Forschung"" through the DLR grant 50.OG.9503.0."
+a suitable normalization must be found.,a suitable normalization must be found.
+ The inlrared is used as à measure of stellar mass to avoid large fluctuations due to a few bright blue stars., The infrared is used as a measure of stellar mass to avoid large fluctuations due to a few bright blue stars.
+ The 2MASS offers a consistent photometric svstem [for (his normalization. however Ferrarese.Coté.&Jordán(2003) ound diserepancies between the X-band magnitudes from the 2\IASS Large Galaxy. Atlas (Jarrettetal.2003) and those derived [rom integrated optical magnitudes and optical to infrared. colors.," The 2MASS offers a consistent photometric system for this normalization, however \citet{fer03} found discrepancies between the K-band magnitudes from the 2MASS Large Galaxy Atlas \citep{jar03} and those derived from integrated optical magnitudes and optical to infrared colors."
+" We chose to use the integrated. extinction corrected. V-band magnitudes yom deVaucouleursetal.(1991). and —Ix), colors from Aaronson(1978). ancl Persson.Matthews.&Aaronson(1973). to perform our normalization."," We chose to use the integrated, extinction corrected V-band magnitudes from \citet{dev91} and $-$ $_0$ colors from \citet{aar78}
+ and \citet{fro78} to perform our normalization."
+ Table G lists the relevant data and references for M81 and the four dwarl ellipticals of this study., Table \ref{tab_nova_rate} lists the relevant data and references for M81 and the four dwarf ellipticals of this study.
+ We included the 24AÀSS Ix-band magnitudes. alter correcting for galactic exGinelion. to compare with the K-band magnitudes derived from the optical magnitudes and colors.," We included the 2MASS K-band magnitudes, after correcting for galactic extinction, to compare with the K-band magnitudes derived from the optical magnitudes and colors."
+ The difference is +0.36 mag for M32 ancl —0.29 mag for NGC 205. illustrating the discrepancy.," The difference is $+$ 0.36 mag for M32 and $-$ 0.29 mag for NGC 205, illustrating the discrepancy."
+" No (V—Ix)j colors were available lor NGC 147 and NGC 135. so we adjusted Koay455,0 by the 0.2 magnitude svstematic offset between the (wo svstems found bv Ferrarese.Coté.&Jordan(2003)."," No $-$ $_0$ colors were available for NGC 147 and NGC 185, so we adjusted $_{2MASS,0}$ by the 0.2 magnitude systematic offset between the two systems found by \citet{fer03}."
+". To arrive at the LSNR. the bulk nova rates are divided by the IX-band huninosities. expressed in 10!L.,κ."," To arrive at the LSNR, the bulk nova rates are divided by the K-band luminosities, expressed in $^{10} L_{\odot,K}$."
+ We plot our results. ancl the result for M81 from Neill&Shara(2004).. as open diamonds in Figure 10. along with data from Ferrarese.Coté.&Jordan(2003).. Table 5. plotted as small filled circles.," We plot our results, and the result for M81 from \citet{nei04}, as open diamonds in Figure \ref{lsnr} along with data from \citet{fer03}, Table 5, plotted as small filled circles."
+ The (wo discrepant points [ον M33 are connected with a dotted line., The two discrepant points for M33 are connected with a dotted line.
+ The upper limits for NGC 147. NGC 185. and NGC 205 are plotted as short horizontal lines with a downward pointing arrow.," The upper limits for NGC 147, NGC 185, and NGC 205 are plotted as short horizontal lines with a downward pointing arrow."
+ The open triangle is (he point [rom (his study for all the dwarl ellipticals assuming the nova candidate in NGC 205 was misclassilied., The open triangle is the point from this study for all the dwarf ellipticals assuming the nova candidate in NGC 205 was misclassified.
+ The horizontal dashed line is (he average luminosity specific nova rate (LSNBR) from (2003)., The horizontal dashed line is the average luminosity specific nova rate (LSNR) from \citet{fer03}.
+. The LSNR for all the dEs that includes the nova candidate in NGC 205 of 177 , The LSNR for all the dEs that includes the nova candidate in NGC 205 of $^{+14.8}_{-4.9}$ 
+"Under these assuniptious. it is possible to derive the following set of jump relations for VCE.B): where is the squared Loreutz factor associated with the normal velocity V, of $.","Under these assumptions, it is possible to derive the following set of jump relations for $\Na(\E,\B)$: where is the squared Lorentz factor associated with the normal velocity $V_{n}$ of ${\cal S}$."
+" Moreover. the discoutinuity of the normal component j, of the current aud that of the charge density i have to satisfy the constraint Quite certainly. Eqs. (10))-(15))"," Moreover, the discontinuity of the normal component $j_{n}$ of the current and that of the charge density $\mu$ have to satisfy the constraint Quite certainly, Eqs. \ref{DiscNasE}) \ref{DiscdnBs}) )"
+" and (17)) have already been established aud used m many other contests. but we are not aware of auv references. and then we give a proof of thei in Appendix A. We are now ready to write the general relation expressing the discoutinuity of the normal derivative of the scalar product E-B asa fiction ofV, aud of the discontiuuities of1 aud p."," and \ref{RelDiscjmu}) ) have already been established and used in many other contexts, but we are not aware of any references, and then we give a proof of them in Appendix A. We are now ready to write the general relation expressing the discontinuity of the normal derivative of the scalar product $\E\cdot\B$ asa function of$V_{n}$ and of the discontinuities of $\j$ and $\mu$."
+" Using equations (12))-(15)). we obtain We now assume that S is a Εξω, ie. there is a charec-separated force-free plasina on one side of S S. sav and a vacuum on the other side."," Using equations \ref{DiscdnEn}) \ref{DiscdnBs}) ), we obtain We now assume that ${\cal S}$ is a FFS, i.e., there is a charge-separated force-free plasma on one side of ${\cal S}$ – ${\cal S}^{-}$, say – and a vacuum on the other side."
+ We thus have the acdditionnal bulk relations iuside the plasima. aud inside the vacuum.," We thus have the additionnal bulk relations inside the plasma, and inside the vacuum."
+ This implies on S. where we have noted to get the last relation that. as a consequence of Eqs. (17))," This implies on ${\cal S}$, where we have noted to get the last relation that, as a consequence of Eqs. \ref{RelDiscjmu}) )"
+" 22d (21))-(25)). 6, o this equality just expressing the fact that S stays a plasina-vactuun interface at anv finie."," and \ref{DiscmuFFS}) \ref{DiscjsFFS}) ), $v_{n}^{-}= V_{n}$ – this equality just expressing the fact that ${\cal S}$ stays a plasma-vacuum interface at any time."
+ Ilereafter. we denote as ( an arcleneth along a magnetic lue. with ( increasing when passing from the plasma to the vactuun. aud as à the associated tangent unit vector (then @-f>0 on 8).," Hereafter, we denote as $\ell$ an arclength along a magnetic line, with $\ell$ increasing when passing from the plasma to the vacuum, and as $\uu$ the associated tangent unit vector (then $\uu\cdot\n\geq 0$ on ${\cal S}$ )."
+ Our aim is to compute the value of the derivative with respect to ( of the component of the electric field., Our aim is to compute the value of the derivative with respect to $\ell$ of the component of the electric field.
+ As E-B=0 ou the plasma side. we have on & Then E satisfics on $ and where we have used equations (29))-(30)) to get the last equality.," As $\E\cdot\B=0$ on the plasma side, we have on ${\cal S}$ Then $E_{\parallel}$ satisfies on ${\cal S}$ and where we have used equations \ref{FFS-dsEB}) \ref{FFS-dnEB}) ) to get the last equality."
+ Using Eqs. (21))-(26)), Using Eqs. \ref{DiscmuFFS}) \ref{EsPlFFS}) )
+" iu the right-haucl side of equation (19)) aud remembering that ο= V,. we get where By using the latter relation iu equation (32)). we obtain finally or. by introducing the angle à between B aud fi. On a part of S not taugeut to B (cosaz 0). we thus have"," in the right-hand side of equation \ref{DiscdnE.B}) ) and remembering that $v_{n}^{-}=V_{n}$ , we get where By using the latter relation in equation \ref{dlEpar}) ), we obtain finally or, by introducing the angle $\alpha$ between $\B$ and $\n$ , On a part of ${\cal S}$ not tangent to $\B$ $\cos\alpha\neq 0$ ), we thus have"
+enerev photons can be scattered so significantly that they become trapped in the svstem.,energy photons can be scattered so significantly that they become trapped in the system.
+ Depending ou their energy. the exact wav in which they escape strongly involves econpetrical effects.," Depending on their energy, the exact way in which they escape strongly involves geometrical effects."
+ Iu the code. we use the escape rate ro“ (or escape probability defined as pos=rs«Rye) derived by Liehtiuaun&Zdziarski(LO87):: where TO is the averaged escape tine. is the Compton interaction probability of photons (of frequency v) with the lepton distributions aud is a relativistic factor correcting for the fact that forward collisions are less effiieut iu trapping the photons.," In the code, we use the escape rate $r_\nu^{\rm{esc}}$ (or escape probability defined as $p^{\rm esc}_\nu=r^{\rm esc}_\nu \times R/c$ ) derived by \citet{LZ87}: where $T_\nu^{\rm{esc}}$ is the averaged escape time, is the Compton interaction probability of photons (of frequency $\nu$ ) with the lepton distributions and is a relativistic factor correcting for the fact that forward collisions are less efficient in trapping the photons."
+ The choice of escape probability is important and different laws can lead to substantially differcutεκ”., The choice of escape probability is important and different laws can lead to substantially different.
+ Although it was shown hat this escape probability reproduces well the results of Monte Carlo simulations in a spherical ecometry (Lightmanetal.nan&Zdziarski 1987).. the conclusions of Sternetal.(1995) nuplv that the escape probability may be shelthy uderestimated.," Although it was shown that this escape probability reproduces well the results of Monte Carlo simulations in a spherical geometry \citep{LZR87,LZ87}, the conclusions of \citet{Stern95} imply that the escape probability may be slightly underestimated."
+ Since the escape Iuniünosity must equal he injected power in a steady state. a sinaller escape xobabilitv miplies a stronger radiation field inside the source.," Since the escape luminosity must equal the injected power in a steady state, a smaller escape probability implies a stronger radiation field inside the source."
+ The consequences on the shape of the photon and lepton distributions become sieuificant onlv when vay production and aunililation are extremely efficient (typically at high optical depths)., The consequences on the shape of the photon and lepton distributions become significant only when pair production and annihilation are extremely efficient (typically at high optical depths).
+ Figure 1 shows spectra in these cases when the deviation from the Moute Carlo simulations become significant., Figure \ref{Coppi92a} shows spectra in these cases when the deviation from the Monte Carlo simulations become significant.
+ Other escape probabilities were proposed (e.g.Sternetal.1995) to reproduce more successfullv the results frou: MC siuulatious iu sole specific regunes. but none were shown to be fully consistent with MC results.," Other escape probabilities were proposed \citep[e.g.][]{Stern95} to reproduce more successfully the results from MC simulations in some specific regimes, but none were shown to be fully consistent with MC results."
+ The use of au escape probability to iunüc ecometrical effects ds of course the iain limitation of our code., The use of an escape probability to mimic geometrical effects is of course the main limitation of our code.
+ Iowever. siguificaut deviations appear only iu optically thick plasmas when steep eradicuts i temperature and intensity. appear. whereas jets and corouae in NRB aud ACN are optically thin media. the largest optical depths observed being τ52. 3.," However, significant deviations appear only in optically thick plasmas when steep gradients in temperature and intensity appear, whereas jets and coronae in XRB and AGN are optically thin media, the largest optical depths observed being $\tau \approx 2-3$ ."
+ The precise geometry of tle sources is also uukuowu and describing accurately the escape probability iu one peculiar ecometry is therefore not necessarily helpful., The precise geometry of the sources is also unknown and describing accurately the escape probability in one peculiar geometry is therefore not necessarily helpful.
+ Particle heating and acceleration are probably amongst of the most mysterious problems of Ligh energy sources., Particle heating and acceleration are probably amongst of the most mysterious problems of high energy sources.
+ Observations show evidence for hot plasmas or uch energv tails in the particle distributions. but little lsi shown about the xecise mechanisms that eenerate these »»pulatious.," Observations show evidence for hot plasmas or high energy tails in the particle distributions, but little is known about the precise mechanisms that generate these populations."
+ Most previous work did. not address this xoblei directly., Most previous work did not address this problem directly.
+ Non-thermal high energy. particles were instead injected iuto the system with an arbitrary (usually oower law) distribution., Non-thermal high energy particles were instead injected into the system with an arbitrary (usually power law) distribution.
+ This ad hoc injection assumes an instantancous acceleration of particles., This ad hoc injection assumes an instantaneous acceleration of particles.
+ It does not ake iuto accom the fact that particle acceleration ju to compete with other cooling processes., It does not take into account the fact that particle acceleration has to compete with other cooling processes.
+ Another «Πίο approach. often used to account for lepton ieatius. consists of assuming that power is provided * sone Unspecified process to the supposedly thermal distribution of clectrous.," Another simple approach, often used to account for lepton heating, consists of assuming that power is provided by some unspecified process to the supposedly thermal distribution of electrons."
+ These prescriptions for particle acceleration aud heating are implemented in the code., These prescriptions for particle acceleration and heating are implemented in the code.
+ Towever. in addition. we also attempted to follow a more plysical approach by implementing two additional specific incchanisius for leating and acceleration. niancly Coulomb heating aud second order Fermi acceleration.," However, in addition, we also attempted to follow a more physical approach by implementing two additional specific mechanisms for heating and acceleration, namely Coulomb heating and second order Fermi acceleration."
+ e c-p Couloiib-like heating: As has already been discussed. ceollisious. with hot xotons cau heat the pair distributions.," $\bullet$ e-p Coulomb-like heating: As has already been discussed, collisions with hot protons can heat the pair distributions."
+ The wav in which ie interaction is adapted iuto the code is described in Sect. 2.1.2., The way in which the interaction is adapted into the code is described in Sect. \ref{e-p}.
+ When the Thomson optical depth is waver than unity. this heating is known to become oeieficicnt and other processes nust operate. which are rot fully understood.," When the Thomson optical depth is lower than unity, this heating is known to become inefficient and other processes must operate, which are not fully understood."
+ À possible means of accounutiug or this additional heating is to mimic the heating by retinal protons but with chhanced efficiency (Navakshin&Melia 1998)., A possible means of accounting for this additional heating is to mimic the heating by thermal protons but with enhanced efficiency \citep{NM98}.
+. Although we do not aiii to describe aux xecise ucroplivsics. this heating prescription estimates cousisteutlv voth FP coefficieuts: the heating rate aud its related diffusion coefficient.," Although we do not aim to describe any precise microphysics, this heating prescription estimates consistently both FP coefficients: the heating rate and its related diffusion coefficient."
+ For this heating prescription. he temperature is set and the total nuniber of rotons is constrained by the initial neutrality.," For this heating prescription, the temperature is set and the total number of protons is constrained by the initial neutrality."
+ The usual cooling and diffusion coefficients 445 (Eq. 22)), The usual cooling and diffusion coefficients $A_{e^\pm}^{e-p}$ (Eq. \ref{eq_A_coul}) )
+ and D. (Eq. 23)), and $D_{e^\pm}^{e-p}$ (Eq. \ref{eq_D_coul}) )
+ are then multiplied bv au efficiency factor jj., are then multiplied by an efficiency factor $\eta$.
+ This efficiency is computed at cach time step so that the total heatingc» is controlled by a coustaut heatingc» parameter order Fermii-like acceleration: This type of acceleration could be genucrated for example by the iuteraction between the electron aud wave turbulence., This efficiency is computed at each time step so that the total heating is controlled by a constant heating parameter $\bullet$ 2nd order Fermi-like acceleration: This type of acceleration could be generated for example by the interaction between the electron and wave turbulence.
+ Diffusion of particles iu mnonieutuni space is then described by the ecucral equation: where f(p) ds the phase-space deusitv., Diffusion of particles in momentum space is then described by the general equation: where $f(p)$ is the phase-space density.
+" When considering an equation about JN, it vields a", When considering an equation about $N_{e^\pm}$ it yields a
+"]€5,I10*—10! are enough for this multi-hand emission. while the turbulent. magnetic field is strong. at least LOC. which is larger than the equipartition value of 1004/6 estimated by svnehrotron radiation (INataokaetal.2006).","$1 \leq \gamma_e \leq
+10^3-10^4$ are enough for this multi-band emission, while the turbulent magnetic field is strong, at least $10^{-3}G$, which is larger than the equipartition value of $100\mu G$ estimated by synchrotron radiation \citep{kataoka06}."
+. The radiative cooling of svnchrotron emission may be one of (he reasons to explain the deeper spectrum toward high energv bands (Ileavens&Aleisenheimerelal. 1989).," The radiative cooling of synchrotron emission may be one of the reasons to explain the deeper spectrum toward high energy bands \citep{heavens87, meisenheimer89}."
+. The observation of supports this traditional interpretation (llarrisetal.—2006)., The observation of supports this traditional interpretation \citep{harris06}.
+. And the svnehrotron enission by (wo population of electrons is needed (Sambrunaetal.2001) (o explain the X-ray spectrum of3C273., And the synchrotron emission by two population of electrons is needed \citep{sambruna01} to explain the X-ray spectrum of.
+. But for the spectrum of the knot inA.. the difference of the spectral indexes between the flatter part and the deeper part is less than 0.5 (Iardcastle. 2006).," But for the spectrum of the knot in, the difference of the spectral indexes between the flatter part and the deeper part is less than 0.5 \citep{hardcastle06}."
+. This is contradictory to the prediction of typical svnchrotron electron cooling., This is contradictory to the prediction of typical synchrotron electron cooling.
+ For another point of view. the relatively low number densitv of the knot can contribute just a small amount of absorption. thus the strong decrease of flux in X-ray band is not due to dust attenuation.," For another point of view, the relatively low number density of the knot can contribute just a small amount of absorption, thus the strong decrease of flux in X-ray band is not due to dust attenuation."
+ Therefore. the cdrop-olf point in the spectrum mieht present the behavior of the turbulence.," Therefore, the drop-off point in the spectrum might present the behavior of the turbulence."
+ From another side. we may directly describe the magnetic field as: D?(&)x&P.," From another side, we may directly describe the magnetic field as: $B^2(k)\propto k^{-p}$."
+ For this point. we avoid the detailed treatments of any turbulence model.," For this point, we avoid the detailed treatments of any turbulence model."
+ With the double power-law as (he form of magnetic field to caleulate DSR. we select py=1.4 and po=1.7 respectively to get the result shown in the Fig.," With the double power-law as the form of magnetic field to calculate DSR, we select $p_1=1.4$ and $p_2=1.7$ respectively to get the result shown in the Fig."
+ 2., 2.
+ But the bulk Lorenz factor is changed from Γι=12 to Py=2., But the bulk Lorenz factor is changed from $\Gamma_1=12$ to $\Gamma_2=2$.
+ This result gives us an alternative clue to explain the multi-band spectrum: the break point in the spectrum might indicate the bulk iransition state of the shock [rom exira-relativistic to sub-relativistic phase., This result gives us an alternative clue to explain the multi-band spectrum: the break point in the spectrum might indicate the bulk transition state of the shock from extra-relativistic to sub-relativistic phase.
+ There are some knots in other objects observed by multi-band. telescopes., There are some knots in other objects observed by multi-band telescopes.
+ Different knots have different. spectral slopes ancl different. quantiGes of flux. indicating the turbulent mode ancl different. acceleration processes.," Different knots have different spectral slopes and different quantities of flux, indicating the non-uniform turbulent mode and different acceleration processes."
+ In (his paper. we give the example ofA.," In this paper, we give the example of."
+. However. whatever the spectral shape is. we see that the observational spectrum can be explained by perturbative DSR. theory. the emission is dominated by the random magnetic field which could be amplified by (he turbulence.," However, whatever the spectral shape is, we see that the observational spectrum can be explained by perturbative DSR theory, the emission is dominated by the random magnetic field which could be amplified by the turbulence."
+ Thus. the spectral shape is uniquely determined by the random magnetic field [rom radio to X-ray band.," Thus, the spectral shape is uniquely determined by the random magnetic field from radio to X-ray band."
+ In (his paper. we use the turbulent spectrum to amplify Che random magnetic field.," In this paper, we use the turbulent spectrum to amplify the random magnetic field."
+ We lind (hat the spectrum shape of DSR. is only dominated by the stochastic magnetic field., We find that the spectrum shape of DSR is only dominated by the stochastic magnetic field.
+ The existence of this kind of magnetic field has been confirmed bv numerical simulations, The existence of this kind of magnetic field has been confirmed by numerical simulations
+local radiation field and the metallicitv.,local radiation field and the metallicity.
+ However. the relative intensities of these PAIL bands appear to be strikingly uniform over these variable local conditions.," However, the relative intensities of these PAH bands appear to be strikingly uniform over these variable local conditions."
+" This work is based on observations made with theTelescope. which is operated by (he Jet Propulsion Laboratory, California Institute of Technology. wider a contract with NASA."," This work is based on observations made with the, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under a contract with NASA."
+ Support lor this work was provided by NASA through contract 1321212. issued by JPL/Caltech to JALC. at Macalester College.," Support for this work was provided by NASA through contract 1321212, issued by JPL/Caltech to J.M.C. at Macalester College."
+ RDG was supported in part by NASA through contracts 1256406 and 1215746 issued by JPL/Caltech to the University of Minnesota., RDG was supported in part by NASA through contracts 1256406 and 1215746 issued by JPL/Caltech to the University of Minnesota.
+ This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. uncer contract with the National Aeronautics and Space Acaministration. and NASA’s Astrophysics Data Svstem.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration, and NASA's Astrophysics Data System."
+ This publication has made use of data products from the Two Micron All sky Survey. which is a joint project of the University. of Massachusetts and (he Infrared Processing and Analvsis Center/California Institute of Technology. Παπάος by the National Aeronautics and Space Administration and the National Science Foundation.," This publication has made use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."
+ We would like to acknowledge Daniel A. Dale. J.D. Smith. Thomas Varberg. and theSpitzer Science Center for helpful discussions and support.," We would like to acknowledge Daniel A. Dale, J.D. Smith, Thomas Varberg, and the Science Center for helpful discussions and support."
+ Finally. we thank (he anonymous referee for a careful and insightful report that improved this manuscript.," Finally, we thank the anonymous referee for a careful and insightful report that improved this manuscript."
+The N-rav emission of Sevfert ealaxics is coummonly believed to be produced by Compton scattering of soft photous ou a population of hot clectrous1980).,The X-ray emission of Seyfert galaxies is commonly believed to be produced by Compton scattering of soft photons on a population of hot electrons.
+ Although high euerev observations of Sevfert ealaxies may be conusisteut with both thermal aud nonthermal models (Zdziarski et al. 1991:, Although high energy observations of Seyfert galaxies may be consistent with both thermal and non–thermal models (Zdziarski et al. \cite{zdz94};
+ Malzac et al. 19983) , Malzac et al. \cite{mal98}) )
+the non.detection of Sevferts by. Comptel aud the high cucrey cut-offs indicated by OSSE (Jourdain ct al. 10051:, the non–detection of Seyferts by Comptel and the high energy cut-offs indicated by OSSE (Jourdain et al. \cite{jou92};
+ Maisack ct al.. 1993))," Maisack et al., \cite{mai93}) )"
+ and BeppoSAN (Matt. 1999)) have focused attention on thermal models., and SAX (Matt \cite{mat99}) ) have focused attention on thermal models.
+ In fact it was shown carly ou that asuning thermal equilibriuni aud a plane parallel ecometry for the Comptonizing plasma above an accretion disk. aud ucelecting direct heating of the disk. the average properties of tle N-rav. emission of Sevtert galaxies could be naturally accounted for (ITaardt Maraschi 1991)).," In fact it was shown early on that assuming thermal equilibrium and a plane parallel geometry for the Comptonizing plasma above an accretion disk, and neglecting direct heating of the disk, the average properties of the X-ray emission of Seyfert galaxies could be naturally accounted for (Haardt Maraschi \cite{haa91}) )."
+ The spectral slope is mainly determined by two parameters. the temperature &T; aud optical depth 7 of the scattering electrons. while the cutoff energy is related essentially to AZ).," The spectral slope is mainly determined by two parameters, the temperature $kT_{\rm e}$ and optical depth $\tau$ of the scattering electrons, while the cut–off energy is related essentially to $kT_e$."
+ Thus. simaltancous measurements of the slope of the A-rayv continuend of the cutoff energy are necessary to determine the plivsical parameters of the Couptouizine reelon.," Thus, simultaneous measurements of the slope of the X-ray continuum of the cut–off energy are necessary to determine the physical parameters of the Comptonizing region."
+ Moreover. im a disk plus corona svsteui the Comptonizine regiou and the source of soft photons arecoupled. as the optically thick disk necessarily reprocesses aud τοσατς part of the Coniptonized flux as soft photons which are the seeds for Coniptouization.," Moreover, in a disk plus corona system, the Comptonizing region and the source of soft photons are, as the optically thick disk necessarily reprocesses and reemits part of the Comptonized flux as soft photons which are the seeds for Comptonization."
+ The svstemi must then satisty equilibrium cucrey balance equations. which depend onycometry aud ou the ratio of direct heating of the disk to that of the corona.," The system must then satisfy equilibrium energy balance equations, which depend on and on the ratio of direct heating of the disk to that of the corona."
+" Iu the limiting case of a “passive” disk. the amplification of the Comptonization process. determined by the Compton parameter y=1KT,)llnoc;)τι|7}. is ⋅⋅fixed by ecometiy ouly."," In the limiting case of a ""passive"" disk, the amplification of the Comptonization process, determined by the Compton parameter $\displaystyle y\simeq 4\left(\frac{kT_e}{m_ec^2}\right)\,
+\left[1+4\left(\frac{kT_e}{m_ec^2}\right)\right] \tau (1+\tau)$, is fixed by geometry only."
+ Therefore.. ifp the corona is iu euergv balance. the temperature and optical depth nuust satisfy a relation which cau e computed for differeut. geometries of the disk|corona configuration (e.g.. the review of Sveussou 1996)).," Therefore, if the corona is in energy balance, the temperature and optical depth must satisfy a relation which can be computed for different geometries of the disk+corona configuration (e.g., the review of Svensson \cite{sve96}) )."
+" It is then theoretically possible to constrain the ecometiy of the svsteimi and verity the sclfcousistency of the model. provided that AZ, and 7 are known with sufficient precision."," It is then theoretically possible to constrain the geometry of the system and verify the selfconsistency of the model, provided that $kT_e$ and $\tau$ are known with sufficient precision."
+ The extent of spectral variability. diving luuinosity variations is an additional. in principle powerful. diagnostic tool that cau be used to test existing models. as it provides direct Iusight iuto the wav the emittius particles are heated aud cooled.," The extent of spectral variability during luminosity variations is an additional, in principle powerful, diagnostic tool that can be used to test existing models, as it provides direct insight into the way the emitting particles are heated and cooled."
+ For example. if the plasima is pair dominated. the value of 7 is fixed by the conrpactuess parameter (1e. bv the hpuuinositv for fixed ecometry).," For example, if the plasma is pair dominated, the value of $\tau$ is fixed by the compactness parameter (i.e. by the luminosity for fixed geometry)."
+ This vields a definite relation between spectral variations and intensity. predicting modest spectral changes for large (factor 10 at least) variatious in the intensity.," This yields a definite relation between spectral variations and intensity, predicting modest spectral changes for large (factor 10 at least) variations in the intensity."
+ On the contrary. for low pair density plasmas. significant spectral variations are possible," On the contrary, for low pair density plasmas, significant spectral variations are possible"
+where the luminosity is in units of erg ! + (2).,where the luminosity is in units of erg $^{-1}$ $^{-1}$ \citep{1983ApJ...266..713O}.
+" One can compute the luminosity evolution for any halo that collapses at redshift ο, and undergoes star formation according to Equation (3)).", One can compute the luminosity evolution for any halo that collapses at redshift $z_c$ and undergoes star formation according to Equation \ref{halo-sfr}) ).
+" The luminosity function at redshift ο, (AZ1g.2). 1s now given by where [IN(M.z.z,.) is the number density at redshift + of halos of mass AJ collapsed at redshift =."," The luminosity function at redshift $z$, $\Phi(M_{AB},z)$, is now given by where $N(M,z,z_c)$ is the number density at redshift $z$ of halos of mass $M$ collapsed at redshift $z_c$."
+ We will use Equation (8)) to study effect of overdensity on the luminosity function and to compare the luminosity function in our model with observations in the next section., We will use Equation \ref{lumfn}) ) to study effect of overdensity on the luminosity function and to compare the luminosity function in our model with observations in the next section.
+ Figure 2. shows the globally averaged luminosity function calculated. using our model for different redshifts in comparison with observations presented by ?.., Figure \ref{lf} shows the globally averaged luminosity function calculated using our model for different redshifts in comparison with observations presented by \citet{2006NewAR..50..152B}.
+ We find that our model reproduces the observed luminosity functions at high redshifts reasonably well., We find that our model reproduces the observed luminosity functions at high redshifts reasonably well.
+ In particular. the match at +=6 is remarkably good while the model predicts less number of galaxies than what is observed at 2=7 and 8.," In particular, the match at $z=6$ is remarkably good while the model predicts less number of galaxies than what is observed at $z=7$ and 8."
+" This could indicate that the star-forming efficiency. f; increases with 2. and/or the time-scale of star formation is lower than /4«4, at higher redshifts."," This could indicate that the star-forming efficiency $f_*$ increases with $z$, and/or the time-scale of star formation is lower than $t_{\rm dyn}$ at higher redshifts."
+ The match of the model with the data can be improved by tuning these parameters suitably. however we prefer not to introduce additional freedom in constraining the parameters: rather our focus is to estimate the effect of reionization and feedback on the luminosity function.," The match of the model with the data can be improved by tuning these parameters suitably, however we prefer not to introduce additional freedom in constraining the parameters; rather our focus is to estimate the effect of reionization and feedback on the luminosity function."
+ In our calculation ofluminosities. we do not make any correction for dust.," In our calculation ofluminosities, we do not make any correction for dust."
+ This is partly because of indications from observed very blue UV-continuum. slopes (2222) that. dust extinction in 27 is small.," This is partly because of indications from observed very blue UV-continuum slopes \citep{2010ApJ...708L..69B,2010ApJ...709L..16O,
+ 2010ApJ...719.1250F,2010MNRAS.tmp.1378B} that dust extinction in $z\gtrsim 7$ is small."
+ As discussed in the next section. we exclusively work with luminosity functions at these redshifts.," As discussed in the next section, we exclusively work with luminosity functions at these redshifts."
+ Also. the effect of dust is degenerate with f; to some extent.," Also, the effect of dust is degenerate with $f_*$ to some extent."
+ Therefore. the exclusion of dust extinction does not affect the general results of our calculation.," Therefore, the exclusion of dust extinction does not affect the general results of our calculation."
+ We have mentioned that galaxy luminosity functions provide valuable information regarding reionization., We have mentioned that galaxy luminosity functions provide valuable information regarding reionization.
+ However. observations are carried— out over —relatively small© fields of view.," However, observations are carried out over relatively small fields of view."
+ The bright sources in these fields are typically hosted by high mass haloes., The bright sources in these fields are typically hosted by high mass haloes.
+ Hence. it is most likely that these fields are biased tracers of luminosity function and reionization.," Hence, it is most likely that these fields are biased tracers of luminosity function and reionization."
+ In this section. we extend our model to study reionization within such biased regions and quantify the departure of various quantities from their globally averaged trends.," In this section, we extend our model to study reionization within such biased regions and quantify the departure of various quantities from their globally averaged trends."
+ Reionization in biased regions has been discussed in the literature., Reionization in biased regions has been discussed in the literature.
+ 2. studied the correlation between high redshift galaxy distribution and the neutral Hydrogen 21 em emission by considering reionization in the vicinity of these galaxies., \citet{2007MNRAS.375.1034W} studied the correlation between high redshift galaxy distribution and the neutral Hydrogen 21 cm emission by considering reionization in the vicinity of these galaxies.
+ ? considered the ionization background near high redshift quasars., \citet{2008MNRAS.383..691W} considered the ionization background near high redshift quasars.
+ ? studied effect of reionization around high redshift quasars on the power spectrum of 21 em emission (see also 25. , \citet{2009MNRAS.399.1877G} studied effect of reionization around high redshift quasars on the power spectrum of 21 cm emission (see also \citealt{2007MNRAS.376.1680P}) ).
+The general conclusion of these studies is that overdense regions are ionised earlier., The general conclusion of these studies is that overdense regions are ionised earlier.
+ In this work. we will consider the behaviour of luminosity functions in such regions.," In this work, we will consider the behaviour of luminosity functions in such regions."
+ Overdense regions are characterised by their comoving Lagrangian size /?. and their linearly extrapolated overdensity 2.," Overdense regions are characterised by their comoving Lagrangian size $R$, and their linearly extrapolated overdensity $\delta$."
+ At a given redshift. we can take a scale £ corresponding to an observed field of view (e... WFC3/IR field in HST) and then determine 0 by identifying the presence of a massive object. for example. a quasar or a bright galaxy.," At a given redshift, we can take a scale $R$ corresponding to an observed field of view (e.g., WFC3/IR field in HST) and then determine $\delta$ by identifying the presence of a massive object, for example, a quasar or a bright galaxy."
+ We follow a prescription discussed by ?.., We follow a prescription discussed by \citet{2008MNRAS.385.2175M}.
+" Note that if a galaxy with luminosity £550 1s observed at redshift z, then we can assign a certain mass to the dark matter halo containing the galaxy. say AZ."," Note that if a galaxy with luminosity $L_{1500}$ is observed at redshift $z_g$ then we can assign a certain mass to the dark matter halo containing the galaxy, say $M$ ."
+ The halo mass AJ has to be, The halo mass $M$ has to be
+solution. if the uncertainties e; were realistic.,"solution, if the uncertainties $\sigma_i$ were realistic."
+ The parücle orbits in comoving galactic coordinates are shown in Figures 1. and 2.., The particle orbits in comoving galactic coordinates are shown in Figures \ref{Fig:1} and \ref{Fig:2}.
+ Thev show two tvpes of motion of LMC and MW., They show two types of motion of LMC and MW.
+ Type I in Models 1. 2 and 5. has LAIC initiallv moving up in a clockwise direction in the ez ancl yz projections.," Type I, in Models 1, 2 and 5, has LMC initially moving up in a clockwise direction in the $xz$ and $yz$ projections."
+ Tvpe II. in Models 3 and 4. has LAIC more directly. approaching MW. trom above with a sharp bend at the low redshilt end.," Type II, in Models 3 and 4, has LMC more directly approaching MW from above with a sharp bend at the low redshift end."
+ These (wo general orbit tvpes. wilh considerable variations in the orbits of M33 and ICLO. are identifiable in all solutions I have found with 4?X60.," These two general orbit types, with considerable variations in the orbits of M33 and IC10, are identifiable in all solutions I have found with $\chi^2\la 60$."
+ The two orbit tvpes largely differ bv the motion of MW relative to the center of mass of the mass mocel ihe motion of LMC relative to MW is quite similar. as will be discussed — but there are two possibly significant points to note.," The two orbit types largely differ by the motion of MW relative to the center of mass of the mass model — the motion of LMC relative to MW is quite similar, as will be discussed — but there are two possibly significant points to note."
+ First. parameter adjustments in the walk to lower V for Type HH solutions often end at an abrupt change of orbits and a large increase in V7.," First, parameter adjustments in the walk to lower $\chi^2$ for Type II solutions often end at an abrupt change of orbits and a large increase in $\chi^2$."
+ The walk in Type I seenis to be approaching a smooth minimum that is not «(uite reached in acceptable computation Gime., The walk in Type I seems to be approaching a smooth minimum that is not quite reached in acceptable computation time.
+ Second. limited (rials with the initial expansion factor changed from equation (3.1)) to a;=0.2 vielded Type II solutions that are very close to what is found al a;—0.1. but did not vield Tvpe I solutions.," Second, limited trials with the initial expansion factor changed from equation \ref{eq:zi}) ) to $a_i=0.2$ yielded Type II solutions that are very close to what is found at $a_i=0.1$, but did not yield Type I solutions."
+ Models 1: and 2 have the smallest 4? found in this study., Models 1 and 2 have the smallest $\chi^2$ found in this study.
+ They. have similar parameters. and the orbits are similar but distinctly different.," They have similar parameters, and the orbits are similar but distinctly different."
+ Doth are shown to illustrate the variations allowed in similar arrangements of the orbits at different apparently local minima of 7., Both are shown to illustrate the variations allowed in similar arrangements of the orbits at different apparently local minima of $\chi^2$.
+ Alodels 3 and 4 are included (o show the alternative Tvpe II behavior., Models 3 and 4 are included to show the alternative Type II behavior.
+ Models 4 ancl 5 illustrate the effect of eliminating the large external mass. at about the lowest 4? for this purpose.," Models 4 and 5 illustrate the effect of eliminating the large external mass, at about the lowest $\chi^2$ for this purpose."
+ AIL 5 models put e above 220 km '., All 5 models put $v_c$ above 220 km $^{-1}$.
+ This may not be significant. however. because ihe nominal catalog value is larger too. and may have biased the choice of local minima.," This may not be significant, however, because the nominal catalog value is larger too, and may have biased the choice of local minima."
+ Remaining to be done is à search for solutions based on a smaller catalog value of v., Remaining to be done is a search for solutions based on a smaller catalog value of $v_c$.
+ since (he masses are allowed considerable departures from nominal it is a positive result that in Models 1 and 2 the mass of M31 is larger (han MW. as is usually considered likely.," Since the masses are allowed considerable departures from nominal it is a positive result that in Models 1 and 2 the mass of M31 is larger than MW, as is usually considered likely."
+ The sums of masses of M31 and MW ave 3.xLOM. and 4x1057.. a substantial difference largely due to the allowed difference of model redshifts of M31.," The sums of masses of M31 and MW are $3\times 10^{12}m_\odot$ and $4\times 10^{12}m_\odot$, a substantial difference largely due to the allowed difference of model redshifts of M31."
+ In these two models the components of proper motion of ICIO differ from the measurement by less than two times the stated error. which seems acceptable. aud (he other proper motions are closer to the measurements.," In these two models the components of proper motion of IC10 differ from the measurement by less than two times the stated error, which seems acceptable, and the other proper motions are closer to the measurements."
+ The clistance of ICI0 in model 2 is large. bul obscuration complicates this measurement.," The distance of IC10 in model 2 is large, but obscuration complicates this measurement."
+ These models place Madfei well away [rom its catalog position. at d.=1.3 and 0.5 Mpc. but this is an acceptable outcome of the idea that Maffei may represent external mass in the general direction of the nearest large galaxy concentration.," These models place Maffei well away from its catalog position, at $d_\perp=1.3$ and $0.5$ Mpc, but this is an acceptable outcome of the idea that Maffei may represent external mass in the general direction of the nearest large galaxy concentration."
+ In both models 4? is consistent with what is expected from the parameter count., In both models $\chi^2$ is consistent with what is expected from the parameter count.
+ Bul since many of the nominal catalog standard deviations are not much better than guesses the conclusion is that the, But since many of the nominal catalog standard deviations are not much better than guesses the conclusion is that the
+ln this way we determined the minimum and maximunmi mass which could have been created in each time interval between output [rom the simulation.,In this way we determined the minimum and maximum mass which could have been created in each time interval between output from the simulation.
+ In Fig., In Fig.
+ 2. we plot the creation rate for halos within two narrow mass ranges., \ref{fig:nbody_form} we plot the creation rate for halos within two narrow mass ranges.
+ In order to obtain the maximum number of creation events. we have used relative low numbers of particles in. cach group.," In order to obtain the maximum number of creation events, we have used relative low numbers of particles in each group."
+ Data are. plotted [or groups of between 4550 and 100110 particles., Data are plotted for groups of between $45-50$ and $100-110$ particles.
+ These cistributions are compared with the three multiplicity functions plotted in Fig. 2..," These distributions are compared with the three multiplicity functions plotted in Fig. \ref{fig:nbody_form},"
+ converted into creation rates by multiplving hy «δα for halos of mass equivalent to 45 or 100 particles., converted into creation rates by multiplying by $d\delta/dt$ for halos of mass equivalent to 45 or 100 particles.
+ These curves have been normalised to the low redshift data., These curves have been normalised to the low redshift data.
+ ALL of the models reproduce the decrease in. creation events to present day seen in the simulations., All of the models reproduce the decrease in creation events to present day seen in the simulations.
+ As output from. the simulation occured. after approximately equal intervals of time. the high. recshift data. sullers as the intervals contain relatively more creation events.," As output from the simulation occured after approximately equal intervals of time, the high redshift data suffers as the intervals contain relatively more creation events."
+ This means that we cannot precisely follow the build-up of the clumps. and the difference between the maximum and minimum mass which could. have been created in cach bin is increased.," This means that we cannot precisely follow the build-up of the clumps, and the difference between the maximum and minimum mass which could have been created in each bin is increased."
+ This is xwticularly noticable in the OCDAM simulation where halos are created at earlier times ancl we have fewer outputs from he simulation., This is particularly noticable in the OCDM simulation where halos are created at earlier times and we have fewer outputs from the simulation.
+ llowever. there is evidence that the solid line (calculated rom the best fit to the mass function) also fits the creation rate data the best out of the three models plotted.," However, there is evidence that the solid line (calculated from the best fit to the mass function) also fits the creation rate data the best out of the three models plotted."
+ As a rough guide to this. the root mean square value. between he plotted cata points and the model is 3.4 for this curve. compared to 6.7 for standard PS theory. ancl 5.0 for the best it model of Sheth Tormen (1999)..," As a rough guide to this, the root mean square value between the plotted data points and the model is 3.4 for this curve, compared to 6.7 for standard PS theory, and 5.0 for the best fit model of Sheth Tormen \shortcite{sheth}."
+ Note that the form of he creation rate is strongly dependent on the parameter e in Equation 20.. and only weakly dependent on parameter p.," Note that the form of the creation rate is strongly dependent on the parameter $a$ in Equation \ref{eq:sheth}, and only weakly dependent on parameter $p$."
+ This is consistent with the importance of these parameters or the mass function: parameter e controls the position of he high-mass cut-olf. whereas parameter p controls the mass tail ofthe distribution.," This is consistent with the importance of these parameters for the mass function: parameter $a$ controls the position of the high-mass cut-off, whereas parameter $p$ controls the low-mass tail of the distribution."
+ We have demonstrated. a simple method. for linking any mass function to the corresponding distribution of times at which isolated halos of a given mass are created., We have demonstrated a simple method for linking any mass function to the corresponding distribution of times at which isolated halos of a given mass are created.
+ In order to provide this link we adopted the assumption that the time scales of interest. are those over which the mass of every chump can be thought of as monotonically increasing., In order to provide this link we adopted the assumption that the time scales of interest are those over which the mass of every clump can be thought of as monotonically increasing.
+ The prior for the collapse time was estimated using the οΕς model which ties in directly. with PS theory. although the method does not use any of PS theory bevond that of the STLIC model.," The prior for the collapse time was estimated using the STHC model which ties in directly with PS theory, although the method does not use any of PS theory beyond that of the STHC model."
+ We have presented. a new derivation of the link between the collapse time and initial overdensity for this model which explicitly shows that this link is independent of the halo mass and is applicable in any Eriedmann cosmology., We have presented a new derivation of the link between the collapse time and initial overdensity for this model which explicitly shows that this link is independent of the halo mass and is applicable in any Friedmann cosmology.
+ Alultiplving the mass function by a function with no mass dependence ancl proportional to the time derivative of the critical overdensity then provides a joint. density function with the correct behaviour for the creation of a halo in massend time., Multiplying the mass function by a function with no mass dependence and proportional to the time derivative of the critical overdensity then provides a joint density function with the correct behaviour for the creation of a halo in mass time.
+ Integrating over the resulting joint density unction will &ive the correct relative number densities of vos within cilferent mass and time intervals., Integrating over the resulting joint density function will give the correct relative number densities of halos within different mass and time intervals.
+ We have extended the analysis of N-bock simulation results presented in paper E to cover three simulations of he build-up of dark matter within cillerent cosmologica models., We have extended the analysis of N-body simulation results presented in paper I to cover three simulations of the build-up of dark matter within different cosmological models.
+ Rather than using PS theory. we have demonstratec row a [it to the mass function may be converted to give a creation rate.," Rather than using PS theory, we have demonstrated how a fit to the mass function may be converted to give a creation rate."
+ Out of the three functions we have compare o the mass function data. the best fit model for these data when converted to a creation rate also fits the creation rate data the best.," Out of the three functions we have compared to the mass function data, the best fit model for these data when converted to a creation rate also fits the creation rate data the best."
+ This gives us confidence that the formalism oesented here is sound. and should give accurate results in more general situations. in particular non-Caussian mocels.," This gives us confidence that the formalism presented here is sound, and should give accurate results in more general situations, in particular non-Gaussian models."
+ We are grateful for the use of the Hydra. N-body code (Couchman.Vhomas&Pearce1995). kindly. provided: by the Hydra consortium.," We are grateful for the use of the Hydra N-body code \cite{couchman}
+ kindly provided by the Hydra consortium."
+resolution (R=A/ÓAz 1500) aud is not well suited for neasurements of faint coutimmum fluxes.,resolution $R\!=\!\lambda/\delta\lambda\!\simeq\!1500$ ) and is not well suited for measurements of faint continuum fluxes.
+ For the SWSOL uode. in particular. s/n aud detector drifts between dark current measurements do not allow measurements of continu enuüssion at levels much below 1 Jv.," For the SWS01 mode, in particular, s/n and detector drifts between dark current measurements do not allow measurements of continuum emission at levels much below 1 Jy."
+" The neasured TRAS 12 and 25 ""coufiuuun fluxes from the SNR would correspoud to ouly 20.01 Jv in the SWS aperture and cannot therefore )o detected with the spectral observations preseuted here.", The measured IRAS 12 and 25 “continuum” fluxes from the SNR would correspond to only $\simeq$ 0.01 Jy in the SWS aperture and cannot therefore be detected with the spectral observations presented here.
+ The spectral sectious including well detected lines are displaved in Fig., The spectral sections including well detected lines are displayed in Fig.
+" 2 and the derived line fluxes are listed in Table 1 together with their contribution to the IRAS fluxes. estimated using the spectral response shown iu Fig. ὃν,"," \ref{sws01} and the derived line fluxes are listed in Table \ref{tab_flux} together with their contribution to the IRAS fluxes, estimated using the spectral response shown in Fig. \ref{iras_filt}."
+ The liue profiles are not resolved within the noise (ic. «OO kin/s). and their centroids are not sienificautly red/blueshifted.," The line profiles are not resolved within the noise (i.e. $<$ 400 km/s), and their centroids are not significantly red/blue–shifted."
+ The most striking result is that the surface brightucss of [NeTIJALL2.8 alone islarger than that observed by IRAS through the 12 filter., The most striking result is that the surface brightness of 12.8 alone is than that observed by IRAS through the 12 filter.
+ This apparent contradiction reflects tle somewhat hijc spatial resolution of ISOSWS xelative to. the a-erssuupled IRAS maps., This apparent contradiction reflects the somewhat higher spatial resolution of ISO–SWS relative to the undersampled IRAS maps.
+ More specifically. the line cluitting filaments visible in the alc [SII] images of Mufsou et al. (1986))," More specifically, the line emitting filaments visible in the and [SII] images of Mufson et al. \cite{mufson}) )"
+ are quite uniformly distributed over παν arcniün. do. al area lareer than the IRAS map resolution. aud the average line πιtLace brightuess over this area is roughly a factor of 21 lower than that within the ISOSWS aperture.," are quite uniformly distributed over many arcmin, i.e. an area larger than the IRAS map resolution, and the average line surface brightness over this area is roughly a factor of 2–4 lower than that within the ISO–SWS aperture."
+ In other words. the ISO spectra. although ceutered on a bright optical filament. does not sample a spot of exceptionally large line surface briehtuess. but rather a “typical” line cluitting reelon.," In other words, the ISO spectra, although centered on a bright optical filament, does not sample a spot of exceptionally large line surface brightness, but rather a “typical” line emitting region."
+ This indicates. therefore. that the IRAS 12 ~coutimmun” is indecd dominated by |NOII| line eiission. i.c. that this line accounts for at least of the observed TRAS flux.," This indicates, therefore, that the IRAS 12 “continuum” is indeed dominated by [NeII] line emission, i.e. that this line accounts for at least of the observed IRAS flux."
+" A similar reasoning applies to the IRAS 25 flux which is strongly coutaminated by enmissiou iu the A226.0. ground state trausition plus significaut contribution fron, [OIV[A225.9.— |SIH]AT1s.7 aud A117.9 (cf."," A similar reasoning applies to the IRAS 25 flux which is strongly contaminated by emission in the 26.0 ground state transition plus significant contribution from 25.9, 18.7 and 17.9 (cf."
+ Table 1). Another interesting result ls that A226.0/|NOIT|AT12.8 and A226.0/A33Ls are both a factor of z1.7 lower than those observed iu RCWIO03 (Oliva ct al. 1998)).," Table \ref{tab_flux}) Another interesting result is that 12.8 and 34.8 are both a factor of $\simeq$ 1.7 lower than those observed in RCW103 (Oliva et al. \cite{rcw103_iso}) ),"
+ while the deusitv seusitive A117.9/A226.0. ratio is similar in the two objects., while the density sensitive 26.0 ratio is similar in the two objects.
+ This result iudicates that IC£13 has a lower Fe gas phase abundance than RCW103 where the Fe/Ne aud Fe/Si relative abundances were fouud to be close to their solar values., This result indicates that IC443 has a lower Fe gas phase abundance than RCW103 where the Fe/Ne and Fe/Si relative abundances were found to be close to their solar values.
+ This difference may reflect intrinsic differences in the ISAL total clement abundances or. more likely. inply that the shock iu IC113 is slower aud herefore less effective in destroving the Febearing grams.," This difference may reflect intrinsic differences in the ISM total element abundances or, more likely, imply that the shock in IC443 is slower and therefore less effective in destroying the Fe–bearing grains."
+ This last possibility is also supported by other “specdommeters”. ie. the line surface brightuess and the [NeII[|/|NoIT| ratio which are lower by a factor of 7 aud 1.6. respectively.," This last possibility is also supported by other “speedometers”, i.e. the line surface brightness and the [NeIII]/[NeII] ratio which are lower by a factor of 7 and 1.6, respectively."
+Z=V.02 modeltheyieldincreasesa fa factoro flwoowinglothe factthallhetemperaluresinthisintermediale— massmodelarehigherthanintheolhermodelsandhencethee[ [ectofourlowerestimale forthe’ F((o.p)? Ne rate is more important.,"$Z=$ 0.02 model the yield increases of a factor of two owing to the fact that the temperatures in this intermediate-mass model are higher than in the other models and hence the effect of our lower estimate for the $\alpha,p)^{22}$ Ne rate is more important."
+ The overall uncertainties in the pproduction due to the uncertainties in the reaction rates are about in the stellar models with mass ~ 3AL... and about in stellar models of lower mass.," The overall uncertainties in the production due to the uncertainties in the reaction rates are about in the stellar models with mass $\simeq$ 3, and about in stellar models of lower mass."
+ For the 5 Z—0.02slelleirmodellheuncertaintiesareabouta [actorof5.duetolhelargeuncertaintieso[ Ihet? F((ap)? Ne rale.," For the 5 $Z=$ 0.02 stellar model the uncertainties are about a factor of 5, due to the large uncertainties of the $\alpha,p)^{22}$ Ne rate."
+ The F(a. p)?Ne reaction rate also influences the production of fluorine in the winds of Woll-Hayet stars hence models of this (ype of stars should also be revised to test the effect of our revised rate and its uncertainties.," The $\alpha,p)^{22}$ Ne reaction rate also influences the production of fluorine in the winds of Wolf-Rayet stars hence models of this type of stars should also be revised to test the effect of our revised rate and its uncertainties."
+ It is also important to note that our estimated lower limit for the PF((o. p)?Ne rate is about 4 orders of magnitude lower than the ?Ne(a. n)? Mg reaction rate.," It is also important to note that our estimated lower limit for the $\alpha,p$ $^{22}$ Ne rate is about 4 orders of magnitude lower than the $^{22}$ $\alpha,n$ $^{25}$ Mg reaction rate."
+" In this case the IE((n.5)?""F reaction has to be taken into account as a possible destruction channel for wwhen a significant neutron flux is released in the convective pulses of AGB stars and in Woll-BRavet stars by the 77 Ne(a. n)?Mg reaction."," In this case the $n,\gamma$ $^{20}$ F reaction has to be taken into account as a possible destruction channel for when a significant neutron flux is released in the convective pulses of AGB stars and in Wolf-Rayet stars by the $^{22}$ $\alpha,n$ $^{25}$ Mg reaction."
+" M the3 Z = 0.02imodelsur [aceabundeancesarealsoshowninFigure llforagivenchoiceoflheparlialinivingzonewil "")02AL...", For the 3 $Z$ =0.02 model surface abundances are also shown in for a given choice of the partial mixing zone with $M_{pmz}$ = 0.002.
+ With the new estimate for the !C(a.5) O rate the contribution of the partial zone is diminished. making this uncertain parameter less important.," With the new estimate for the $^{14}$ $\alpha,\gamma$ $^{18}$ O rate the contribution of the partial mixing zone is diminished, making this uncertain parameter less important."
+ In. particular. in the lower limit case. the resulting //!O]] ratio is the same within as computed without including the partial mixing zone (compare to Figure»5)).," In particular, in the lower limit case, the resulting ] ratio is the same within as computed without including the partial mixing zone (compare to )."
+ In none of the cases we calculated could the highest //!O]]| values observed be reproduced., In none of the cases we calculated could the highest ] values observed be reproduced.
+ As discussed in83.. (his problem should be reviewed with the inclusion in future calculations of extra-mixing processes (Cool Bottom Processing) at the base of the convective envelope.," As discussed in, this problem should be reviewed with the inclusion in future calculations of extra-mixing processes (Cool Bottom Processing) at the base of the convective envelope."
+ Future work should also improve our knowledge of the formation ancl the nucleosyntliesis in (he partial mixing zone., Future work should also improve our knowledge of the formation and the nucleosynthesis in the partial mixing zone.
+ One hypothesis is (hat rotation can play a role in varving the efficienev of the production of aand of the eelements (IIerwigetal.2003)., One hypothesis is that rotation can play a role in varying the efficiency of the production of and of the elements \citep{herwig:03}.
+. It will be of much interest to analvze the effects of this hypothesis on the correlation between fIuorine and the s-process elements and (o revise the available observational data., It will be of much interest to analyze the effects of this hypothesis on the correlation between fluorine and the $s$ -process elements and to revise the available observational data.
+ Using data for carbon stars [roii Utsumi(1985). it appeared that these (wo quantities were correlated in AGB stars. however using more recent and precise data from Abiaetal.(2002) this correlation does not seem to appear anymore.," Using data for carbon stars from \citet{utsumi:85} it appeared that these two quantities were correlated in AGB stars, however using more recent and precise data from \citet{abia:02} this correlation does not seem to appear anymore."
+ Another problem is related to C(J) stars., Another problem is related to C(J) stars.
+ Ht is still unknown if (hese stars actually belong io the AGB group or if they are in some other phase of the evolution., It is still unknown if these stars actually belong to the AGB group or if they are in some other phase of the evolution.
+ Moreover. it appears that their ΟΙ ratios around 0.6 are due to a low abundance of," Moreover, it appears that their ] ratios around 0.6 are due to a low abundance of"
+The United Ixingdom LIafrared Telescope is operated by the Joint Astronomy Centre on behalf of the U.Ix. Science and Technology. Facilities Council.,The United Kingdom Infrared Telescope is operated by the Joint Astronomy Centre on behalf of the U.K. Science and Technology Facilities Council.
+ Partly based on observations carried out with the ESO telescopes under programmes 077.X-0667 ancl 177.N-0591., Partly based on observations carried out with the ESO telescopes under programmes 077.A-0667 and 177.A-0591.
+" Based on. observations obtained at the Giomini Observatory, which is operated. by the Association of Universities for Research in Astronomy. Inc. uncer a cooperative agreement with the NSE on behalf of the Gemini partnership: the National Science. Foundation (United States). the Particle Physics and Astronomy Research Council (United Ixingdom). the National Rescarch Council (Canada). CONICYT. (Chilo). the Australian Research Council (Australia). CNPq (Brazil) and CONICET (Argentina)."," Based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Particle Physics and Astronomy Research Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), CNPq (Brazil) and CONICET (Argentina)."
+ The Faulkes Telescopes are operated. by the Las Cumbres Observatory., The Faulkes Telescopes are operated by the Las Cumbres Observatory.
+ The Dark Cosmology Centre is funded by the Danish National Research Foundation., The Dark Cosmology Centre is funded by the Danish National Research Foundation.
+"In practice. gas is present at a range of excitation temperatures. from ο 1001ν to ~2000Ix. Typical Gvo-temperature fits to the extinction diagrams for ULIRGs*. such as those in Figure GSSupp. indicate that ~107M. of gas is found at Z7Z30018 and ~10"". of gas is found al T= 1000I. The massive colder component contributes very little to lines with high J values. but dominates the [huxes of low-7 lines. while a small amount of warmer gas boosts hieh-J [luxes.","In practice, gas is present at a range of excitation temperatures, from $\sim$ 150K to $\sim$ 2000K. Typical two-temperature fits to the extinction diagrams for $^3$, such as those in Figure \ref{pic_excitation}S Supp, indicate that $\sim 10^{9}M_{\odot}$ of gas is found at $T\la 300$ K and $\sim 10^6M_{\odot}$ of gas is found at $T\ga 1000$ K. The massive colder component contributes very little to lines with high $J$ values, but dominates the fluxes of $J$ lines, while a small amount of warmer gas boosts $J$ fluxes."
+ To quantily deviations Iron a single-temperature spectrum. for any three lines I can compute A244/.J5.Jy]. which is the ratio of the observed flux of the ο.) line to that expected from lines 5(45) and (ο) if all Uavee transitions were from (he same excitation temperature.," To quantify deviations from a single-temperature spectrum, for any three lines I can compute $R[J_1/J_2,J_3]$, which is the ratio of the observed flux of the $J_1$ ) line to that expected from lines $J_2$ ) and $J_3$ ) if all three transitions were from the same excitation temperature."
+ In (he case of mulü-temperature gas al an ecquilibrium ortho-to-para ratio and with no extinction. excitation diagrams are concave and / values are less than 1 (lor Jj between 4» and JJ).," In the case of multi-temperature gas at an equilibrium ortho-to-para ratio and with no extinction, excitation diagrams are concave and $R$ values are less than 1 (for $J_1$ between $J_2$ and $J_3$ )."
+ Bolometric [Iuxes awe caleulated by fitting Ηνο or bbroad-band data: using a single-temperature black body. function modified by emissivity xv! , Bolometric fluxes are calculated by fitting IRAS or broad-band data using a single-temperature black body function modified by emissivity $\propto \nu^{-1}$.
+The values obtained this way for IRAS data are within of those obtained using the fitting formula where [ο ave IRAS [κος in Jv.," The values obtained this way for IRAS data are within of those obtained using the fitting formula where $F_{12, ...}$ are IRAS fluxes in $^1$."
+ The acenracy of PAIL flux measurements is limited by the assumptions that were mace about the shapes of the features which are fixed in my analvsis., The accuracy of PAH flux measurements is limited by the assumptions that were made about the shapes of the features which are fixed in my analysis.
+ However. the results in this Letter appear to be robust to the fitting method used for ealeulating PATI fluxes.," However, the results in this Letter appear to be robust to the fitting method used for calculating PAH fluxes."
+ For example. using PAIT fluxes obtained by sunning up the spectrum above the local linear continuunp in Figures 2a.c produces similar results and the correlations with 59.7/an]| are easily detectable.," For example, using PAH fluxes obtained by summing up the spectrum above the local linear $^{8}$ in Figures 2a,c produces similar results and the correlations with $S[9.7\micron]$ are easily detectable."
+" As for PAIIL[6.2,/m]] (Figure 2b). it needs to be corrected. [or water ice absorption (i.e.. multiplied by exp(5|6.0jan]) before the correlation of ΡΑΠΤΗ/PATE[6.27/m]] with 59.7,0n] becomes apparent."," As for ] (Figure 2b), it needs to be corrected for water ice absorption (i.e., multiplied by $\exp(S[6.0\micron])$ before the correlation of ] with $S[9.7\micron]$ becomes apparent."
+ All conclusions remain the same if the main sample of ULIRGs is split by optical classification (LINERs vs HII regions) ancl each category is considered separately., All conclusions remain the same if the main sample of ULIRGs is split by optical classification (LINERs vs HII regions) and each category is considered separately.
+ Previous studies of mid-IR llines adopted varving views on the ellects of absorption., Previous studies of mid-IR lines adopted varying views on the effects of absorption.
+" For ULIRGs observed with ISO. [fluxes were corrected for absorption belore constructing excitation diagrams""."," For ULIRGs observed with ISO, fluxes were corrected for absorption before constructing excitation $^{36}$."
+" More recently"" zero absorplion correction was suggested [or iin (he auxiliary sample on the basis of simultaneous measurements of S(1)/5(3) aud S(1)/9(2) ratios as a function of the excitation temperature (this argument is analogous to saving (hat R[2/1.3] values in ULIRGs agree with the model value)."," More $^3$ zero absorption correction was suggested for in the auxiliary sample on the basis of simultaneous measurements of S(1)/S(3) and S(1)/S(2) ratios as a function of the excitation temperature (this argument is analogous to saying that $R[2/1,3]$ values in ULIRGs agree with the model value)."
+" It is largely unknown if the ro-vibrational lines in ULIRGs are co-spatial with the mid-I. rotational lines: previous studies of ULIRGs in the near-IR tend to correct the entire spectrum [for extinction"".", It is largely unknown if the near-IR ro-vibrational lines in ULIRGs are co-spatial with the mid-IR rotational lines; previous studies of ULIRGs in the near-IR tend to correct the entire spectrum for $^{37}$ .
+" (5z5). TzLON T=107K IIl» Z4; ~Αι A~0.05 Adqnas eis (Tei.xmLOS) Ίσα ACDML O,,=0.3.0,OLO40.7.0.65. a,=1. O,/0,,. "," $z\gsim 5$ $T\gsim 10^4$ $T\gsim 10^2$ ${\rm H_2}$ $R_{\rm vir}$ $\sim \lambda R_{\rm vir}$ $\lambda\sim
+0.05$ $\sim M_{\rm disk} v_c^2$$M_{\rm disk}$ $v_c$ $T_{\rm vir}\gsim 10^6$ $\alpha$ $\Lambda$ $\Omega_m=0.3,
+\Omega_b=0.04, \Omega_\Lambda=0.7, h=0.65$ $\sigma_8=1$ $\Omega_b/\Omega_m$ "
+Ileve 6=Ed is a dimensionless parameter (hat. enters when we choose to normalize ihe magnete flux &=c/c where ty is a (vpical © value.,Here $\delta =\frac{HB_0}{\psi_0}$ is a dimensionless parameter that enters when we choose to normalize the magnetic flux $\bar{\psi} =\psi /\psi_0$ where $\psi_0$ is a typical $\psi$ value.
+ In practice we only work with dimensionless quantities and drop the bars without confusion., In practice we only work with dimensionless quantities and drop the bars without confusion.
+"incorporated in the models have been described previously (Chaboyeretal.2001;Bjork&Chaboyer2006;Dotteretal.2007, 2008),, but we shall provide a brief summary.","incorporated in the models have been described previously \citep{Chaboyer2001,Bjork2006,Dotter2007,Dotter2008}, but we shall provide a brief summary."
+" Our models include the effects of helium and heavy element diffusion following the prescription of Thouletal.(1994),, though for fully convective stars, diffusion physics are unimportant."," Our models include the effects of helium and heavy element diffusion following the prescription of \citet{Thoul1994}, though for fully convective stars, diffusion physics are unimportant."
+ The opacities utilized by DSEP are the high-temperature OPAL opacities (IglesiasRogers1996) and low-temperature opacities of Fergusonetal.," The opacities utilized by DSEP are the high-temperature OPAL opacities \citep{Iglesias1996}
+ and low-temperature opacities of \citet{Ferguson2005}."
+" Surface boundary conditions were defined using the (2005)..PHOENIX model atmospheres (Hauschildtetal.1999a,b) and were attached to the interior model at T=Tepe by interpolating in model atmosphere tables."," Surface boundary conditions were defined using the PHOENIX model atmospheres \citep{Hauschildt1999a,Hauschildt1999b} and were attached to the interior model at $T=T_{eff}$ by interpolating in model atmosphere tables."
+" Attaching the atmospheres to T=T,ry makes thevalue of the convective mixing length used in the atmosphere code inconsequential (Baraffeetal.", Attaching the atmospheres to $T=T_{eff}$ makes thevalue of the convective mixing length used in the atmosphere code inconsequential \citep{BCAH97}.
+" Above 0.8 Mo, DSEP uses a general 1997)..ideal gas equation of state with the Debye-Hiicckel correction (Chaboyer&Kim 1995)."," Above 0.8 $M_{\odot}$ , DSEP uses a general ideal gas equation of state with the Debye-Hücckel correction \citep{Chaboyer1995}."
+". In the low-mass regime, DSEP employs the in the EOS4 configuration, selected for its treatment of arbitrary heavy element abundances and its inclusion of the H} molecule."," In the low-mass regime, DSEP employs the in the EOS4 configuration, selected for its treatment of arbitrary heavy element abundances and its inclusion of the $_{2}^+$ molecule."
+ FreeEOS has also been shown to be valid for modeling stars more massive than 0.1 Mo (Irwin2007)., FreeEOS has also been shown to be valid for modeling stars more massive than 0.1 $M_{\odot}$ \citep{Irwin2007}.
+. Convective core overshoot is included using the method of Demarqueetal.(2004)., Convective core overshoot is included using the method of \citet{Demarque2004}.
+. Rotation was not considered., Rotation was not considered.
+ The only modification made to the underlying physics in DSEP is related to the partial inhibition of element diffusion., The only modification made to the underlying physics in DSEP is related to the partial inhibition of element diffusion.
+ We have introduced turbulent diffusion as described by Richardetal.(2005)., We have introduced turbulent diffusion as described by \citet{Richard2005}.
+. Turbulent diffusion modifies the atomic diffusion coefficient and acts to extend the mixing region below the convection zone., Turbulent diffusion modifies the atomic diffusion coefficient and acts to extend the mixing region below the convection zone.
+" The magnitude of the turbulent diffusion coefficient is tied to an adjustable reference temperature, 10, and varies with density via: where w characterizes the relative strength of turbulent diffusion and Dy-(To) and p(To) are the helium diffusion coefficient and density at the prescribed reference temperature, respectively."," The magnitude of the turbulent diffusion coefficient is tied to an adjustable reference temperature, $T_{0}$, and varies with density via: where $\omega$ characterizes the relative strength of turbulent diffusion and $D_{He}(T_{0})$ and $\rho(T_{0})$ are the helium diffusion coefficient and density at the prescribed reference temperature, respectively."
+" Proffitt&Michaud(1991) motivate the p-? dependence in order to reproduce the solar beryllium abundance, which appears to be unchanged over time."," \citet{Proffitt1991} motivate the $\rho^{-3}$ dependence in order to reproduce the solar beryllium abundance, which appears to be unchanged over time."
+" Thus, any non-standard mixing in the Sun must be localized to a narrow region below the solar convection zone."," Thus, any non-standard mixing in the Sun must be localized to a narrow region below the solar convection zone."
+" We select w—400 and leave it fixed, in concordance with Richardetal.(2005)."," We select $\omega = 400$ and leave it fixed, in concordance with \citet{Richard2005}."
+" 'The reference temperature used primarily in our models is Το= 109, which was found to best reproduce the observed abundance trends of (Kornetal. 2007)."," The reference temperature used primarily in our models is $T_{0} = 10^{6}$ , which was found to best reproduce the observed abundance trends of \citep{Korn2007}."
+". A clar calibration model was generated to determine the appropriate initial mass fractions of helium (Yinit) and metals (Zinit), given the solar heavy element composition of Grevesse&Sauval(1998),, as well as to calibrate the convective mixing length, (arr=£/Hp)."," A solar calibration model was generated to determine the appropriate initial mass fractions of helium $Y_{init}$ ) and metals $Z_{init}$ ), given the solar heavy element composition of \citet{GS98}, as well as to calibrate the convective mixing length, $\left(\alpha_{MLT} = \ell / H_P\right)$."
+" At the solar age (4.57Gyr;Bahcalletal.2005) we were able to reproduce the solar radius, solar luminosity, radius of the convective boundary, and with OMLT = 1.938, Yinit => 0.27491, and Zinit = (Z/X)o0.01884."," At the solar age \citep[4.57 Gyr;][]{Bahcall2005} we were able to reproduce the solar radius, solar luminosity, radius of the convective boundary, and $(Z/X)_{\odot}$ with $\alpha_{MLT}$ = 1.938, $Y_{init}$ = 0.27491, and $Z_{init}$ = 0.01884."
+ All of the models utilized in this study were calculated using the solar calibration as a reference and were assumed to be coeval., All of the models utilized in this study were calculated using the solar calibration as a reference and were assumed to be coeval.
+ Super-solar metallicity models with [Fe/H] = +0.15 were generated with Yinit = 0.28419 and Zinit = 0.02469., Super-solar metallicity models with [Fe/H] = +0.15 were generated with $Y_{init}$ = 0.28419 and $Z_{init}$ = 0.02469.
+" The age of the system was constrained by evolving a 1.347 Mo model, with [Fe/H] = +0.15 and solar calibrated orr, and matching the model radius with the observed radius of KOI-126 A (Figure 1)."," The age of the system was constrained by evolving a 1.347 $M_{\odot}$ model, with [Fe/H] = +0.15 and solar calibrated $\alpha_{MLT}$, and matching the model radius with the observed radius of KOI-126 A (Figure 1)."
+" The age we derive for the system is 4.1 + 0.6 Gyr, consistent with the age derived by C11."," The age we derive for the system is 4.1 $\pm$ 0.6 Gyr, consistent with the age derived by C11."
+" Our uncertainty in the age is dominated by the uncertainty in the observed mass and metallicity of the system, with the observational uncertainty of the radius being of negligible importance."," Our uncertainty in the age is dominated by the uncertainty in the observed mass and metallicity of the system, with the observational uncertainty of the radius being of negligible importance."
+ The effect of changing the reference temperature for turbulent diffusion was also considered (including removing it entirely) and was found to play a negligible role in the age determination., The effect of changing the reference temperature for turbulent diffusion was also considered (including removing it entirely) and was found to play a negligible role in the age determination.
+ The primary results of this Letter are demonstrated in Figure 2., The primary results of this Letter are demonstrated in Figure 2.
+ DSEP accurately reproduces the observed radius for each of the low-mass stars at 4.1 Gyr with [Fe/H] = +0.15., DSEP accurately reproduces the observed radius for each of the low-mass stars at 4.1 Gyr with [Fe/H] = +0.15.
+" We find that for masses Mz = 0.2410 Mo and Ma = 0.2130 Mo the predicted radii from DSEP are Ro = 0.2544 Ro and R3 = 0.2312 Ro, indicating a relative error between the model and observed radii of less than0."," We find that for masses $_2$ = 0.2410 $M_{\odot}$ and $_3$ = 0.2130 $M_{\odot}$ the predicted radii from DSEP are $_2$ = 0.2544 $R_{\odot}$ and $_3$ = 0.2312 $R_{\odot}$ , indicating a relative error between the model and observed radii of less than."
+"3%.. At solar metallicity, the models predict radii approximately smaller than those predicted by the super-solar metallicity models."," At solar metallicity, the models predict radii approximately smaller than those predicted by the super-solar metallicity models."
+ The predicted radii are robust., The predicted radii are robust.
+ Artificially reducing the mixing length (αι=1.00) and fitting the atmosphere to a deeper point in the stellar envelope (r= 100) both produced radius changes under0.5%.., Artificially reducing the mixing length $\alpha_{MLT}=1.00$) and fitting the atmosphere to a deeper point in the stellar envelope $\tau =100$ ) both produced radius changes under.
+ Solar metallicity isochrones from DSEP display radii approximately larger than radii predicted by BCAH98., Solar metallicity isochrones from DSEP display radii approximately larger than radii predicted by BCAH98.
+ The difference islikely a consequence of the, The difference islikely a consequence of the
+et al.,et al.
+" 2009 provide a much improved algorithm for mapping out the stellar mass density, which makes use of multi-band images."," 2009 provide a much improved algorithm for mapping out the stellar mass density, which makes use of multi-band images."
+" However, here we do not require the exact amplitude of stellar mass density variations, but only the location of mass density enhancements to define thelocation of spiral arms."," However, here we do not require the exact amplitude of stellar mass density variations, but only the location of mass density enhancements to define the of spiral arms."
+ Foreground stars in the processed 3.6m images are removed according to the UV color cut described above and the images are deprojected according to the values in Table 1., Foreground stars in the processed $\mu$ m images are removed according to the UV color cut described above and the images are deprojected according to the values in Table 1.
+" In order to make the 3.6m surface brightness variations a better approximation to thelocation of stellar mass density enhancements, we apply some spatial filtering."," In order to make the $\mu$ m surface brightness variations a better approximation to the of stellar mass density enhancements, we apply some spatial filtering."
+" The 3.6um images are first median filtered over 20 pixels kpc), to remove bright spots and features."," The $\mu$ m images are first median filtered over 20 pixels $\approx$ 1 kpc), to remove bright spots and features."
+ The filtered (+1image is then Fourier-decomposed in $ with radial bins that overlap to obtain a version of the image that consists of the m=6 component divided by the m=0 component., The filtered image is then Fourier-decomposed in $\phi$ with radial bins that overlap to obtain a version of the image that consists of the m=6 component divided by the m=0 component.
+" This spatially filters the images, which are shown in Figure 1.."," This spatially filters the images, which are shown in Figure \ref{mask}."
+" We note that the grand-design spirals, NGC 5194 and NGC 628, have spiral arms that were well-defined using only the m—4 component divided by the m=0 component and that moving up to m=6 leads to little, if any difference."," We note that the grand-design spirals, NGC 5194 and NGC 628, have spiral arms that were well-defined using only the m=4 component divided by the m=0 component and that moving up to m=6 leads to little, if any difference."
+" However, in the more flocculent spiral, NGC 6946, the 3.6um structure is more complex, requiring an extension to m=6."," However, in the more flocculent spiral, NGC 6946, the $\mu$ m structure is more complex, requiring an extension to m=6."
+" The inner bulge regions, where no spiral arms are evident, are masked out in each galaxy."," The inner bulge regions, where no spiral arms are evident, are masked out in each galaxy."
+" The outer limit is set by the area over which the CO maps detect emission, or, in the case of NGC 5194, when the arms become too tightly wound for accurate definition by the mask."," The outer limit is set by the area over which the CO maps detect emission, or, in the case of NGC 5194, when the arms become too tightly wound for accurate definition by the mask."
+ In all three cases the analysis does not extend to the outer regions of the galaxies., In all three cases the analysis does not extend to the outer regions of the galaxies.
+" We extend to 0.3 or 0.4r25 depending on the galaxy, where most of the star formation and luminosity is found."," We extend to 0.3 or $r_{25}$ depending on the galaxy, where most of the star formation and luminosity is found."
+ The inner and outer limits for each galaxy are listed in Table 1., The inner and outer limits for each galaxy are listed in Table 1.
+" The image is then divided into radial annuli each of 7.5"" width.", The image is then divided into radial annuli each of $''$ width.
+" We chose a width below the resolution of the image in order to ensure overlapping radial bins, which produce a more continuous spiral arm structure in the masks."," We chose a width below the resolution of the image in order to ensure overlapping radial bins, which produce a more continuous spiral arm structure in the masks."
+" In each annulus, the ‘arm region’ is defined to be the area covered by a certain percentage of the highest-value pixels (e.g., the brightest 30%))."," In each annulus, the `arm region' is defined to be the area covered by a certain percentage of the highest-value pixels (e.g., the brightest )."
+" In a similar fashion, the ‘interarm region’ is defined as the area covered by the same percentage of the lowest-value pixels."," In a similar fashion, the `interarm region' is defined as the area covered by the same percentage of the lowest-value pixels."
+ We will vary the exact percentage used to define these regions over the range tto ensure conclusions robust to the precise arm definition., We will vary the exact percentage used to define these regions over the range to ensure conclusions robust to the precise arm definition.
+ We refer to the percentage used to define the arms as the ‘arm pixel fraction’., We refer to the percentage used to define the arms as the `arm pixel fraction'.
+" Once both the arm and interarm regions each consist of of the pixels, the entire surface area is covered."," Once both the arm and interarm regions each consist of of the pixels, the entire surface area is covered."
+ Figure 1 shows the 3.6um image (far left) and the Fourier reconstructed m=6 image divided by the m=0 (second from left) with contours overlaid showing the arm regions , Figure \ref{mask} shows the $\mu$ m image (far left) and the Fourier reconstructed m=6 image divided by the m=0 (second from left) with contours overlaid showing the arm regions (white).
+The spiral arm masks from right) for our (white).sample where the arm regions are (seconddefined from the highest pixels per radial bin are also shown as is the 24um image with arm contours overlaid for comparison., The spiral arm masks (second from right) for our sample where the arm regions are defined from the highest pixels per radial bin are also shown as is the $\mu$ m image with arm contours overlaid for comparison.
+" Having defined the arm and interarm regions, we now can assess what fraction of star formation tracers is found in the respective regions."," Having defined the arm and interarm regions, we now can assess what fraction of star formation tracers is found in the respective regions."
+" To do so, we measure the fraction of the total emission from gas and star formation tracers that are contained in the arm regions for the area considered (see Table 1 for inner and outer radii)."," To do so, we measure the fraction of the total emission from gas and star formation tracers that are contained in the arm regions for the area considered (see Table 1 for inner and outer radii)."
+ We make this measurement for a variety of arm pixel fractions., We make this measurement for a variety of arm pixel fractions.
+" We examine the HI, H2 (as traced by CO), total gas, FUV, 24 wm, SFR, stellar mass surface density."," We examine the HI, $_{2}$ (as traced by CO), total gas, FUV, 24 $\mu$ m, SFR, stellar mass surface density."
+" The plots in Figure 2 show the fraction of the overall emission found in the arm regions, as a function of the arm pixel fraction."," The plots in Figure \ref{six} show the fraction of the overall emission found in the arm regions, as a function of the arm pixel fraction."
+ Figure 2 a) shows the flux fraction for different tracers that occur in the arm region as a function of the pixel fraction used to define the arm region., Figure \ref{six} a) shows the flux fraction for different tracers that occur in the arm region as a function of the pixel fraction used to define the arm region.
+" To emphasize the difference between the curves, we have divided by the expectation of a spatially uniform distribution and shown them in b)."," To emphasize the difference between the curves, we have divided by the expectation of a spatially uniform distribution and shown them in b)."
+ The 3.6jm image is included in all panels for ease of comparison., The $\mu$ m image is included in all panels for ease of comparison.
+ The thin black line in a) denotes what would be expected for an azimuthally uniform distribution., The thin black line in a) denotes what would be expected for an azimuthally uniform distribution.
+" In b), an azimuthally uniform tracer would be a horizontal line (flux enhancement — 1)."," In b), an azimuthally uniform tracer would be a horizontal line (flux enhancement = 1)."
+ For ease of comparison in our discussion we use a fiducial pixel fraction of to define the arms., For ease of comparison in our discussion we use a fiducial pixel fraction of to define the arms.
+ Figure 2 shows that all tracers are much more concentrated to the spiral arms than a uniform distribution (i.e. all values are above unity in panel b))., Figure \ref{six} shows that all tracers are much more concentrated to the spiral arms than a uniform distribution (i.e. all values are above unity in panel b)).
+" For the 3.6u4m image this is by construction, as the wavelength was used to define the arm regions."," For the $\mu$ m image this is by construction, as the wavelength was used to define the arm regions."
+" However, Figure 2 shows that all the tracers of star formation"," However, Figure \ref{six} shows that all the tracers of star formation"
+The supernova models were evolved with to the point where all explosive unclear burning bad ceased aud the reverse shock had just beguu to form.,The supernova models were evolved with to the point where all explosive nuclear burning had ceased and the reverse shock had just begun to form.
+ This occurred at 107 s and 105 s for 15 and 25 sstars of solar composition. and at 25 s and 100 8 for 15 and 25 sstars of primordial composition. respectively.," This occurred at $10^3$ s and $10^4$ s for 15 and 25 stars of solar composition, and at $25$ s and $100$ s for 15 and 25 stars of primordial composition, respectively."
+ At these nues. the one dinensioual models frou: were uapped onto a two-dimensional axisvuunetric erid and evolved forward in time with the code.," At these times, the one dimensional models from were mapped onto a two-dimensional axisymmetric grid and evolved forward in time with the code."
+ A similar simulation of Model s25A was also performed using a xoseuitor evolved to 10° secouds with before cine mapped toFLASH., A similar simulation of Model s25A was also performed using a progenitor evolved to $10^3$ seconds with before being mapped to.
+. No significant difference in the ater evolution of 325A models evolved to 107 seconds aud 104 seconds was observed., No significant difference in the later evolution of s25A models evolved to $10^3$ seconds and $10^4$ seconds was observed.
+ Oulv one quadrant of the star was carried in the calculation. cuforciug sviunetry about he left y- aud bottom «-axes. while allowing material to cave the erid through a zero-gradieut bouudary at the right y- aud top .c-axes.," Only one quadrant of the star was carried in the calculation, enforcing symmetry about the left $y$ - and bottom $x$ -axes, while allowing material to leave the grid through a zero-gradient boundary at the right $y$ - and top $x$ -axes."
+ An chhauced flow was observed along the ο and g-, An enhanced flow was observed along the $x$ - and $y$ -axes.
+ This flow was not large iu comparison with the vest of the simulation. but if was preseut. as can be seen in Figures 7. - 6.," This flow was not large in comparison with the rest of the simulation, but it was present, as can be seen in Figures \ref{s15A_dens} - \ref{z25D_dens}."
+ This is a welbdocunenuted artifact of the dimensionally-split approach to solving the lvdvodvuamic equations., This is a well-documented artifact of the dimensionally-split approach to solving the hydrodynamic equations.
+ It did not substantially influence the evolution of the simulations preseuted here., It did not substantially influence the evolution of the simulations presented here.
+ Perturbations arising from a Cartesian eril are also inevitable., Perturbations arising from a Cartesian grid are also inevitable.
+ Iu order to quantify these grid effects. we performed smimlatious of all stars with a random perturbation in velocity with a maximum amplitude of 0.54 and 2.0%.," In order to quantify these grid effects, we performed simulations of all stars with a random perturbation in velocity with a maximum amplitude of $0.5\%$ and $2.0\%$."
+ We also performed simmlations with uo additional perturbation., We also performed simulations with no additional perturbation.
+ We fouud that perturbations of 2.0% in velocity had a clear effect on the initial scale of the Ravleigh-Tavlor instability. increasing the auplitude and scale of the first instabilities to form.," We found that perturbations of $2.0\%$ in velocity had a clear effect on the initial scale of the Rayleigh-Taylor instability, increasing the amplitude and scale of the first instabilities to form."
+ This effect is shown for sL5A in Figure L.., This effect is shown for s15A in Figure \ref{s15A_pert}.
+ The case of 0.5% raudom perturbations to the velocity results in a scale for the iuitial Ravleigh-Tavlor instabilities that is roughly equivalent to the case where the ouly perturbations were those arising from the Cartesian grid itself., The case of $0.5\%$ random perturbations to the velocity results in a scale for the initial Rayleigh-Taylor instabilities that is roughly equivalent to the case where the only perturbations were those arising from the Cartesian grid itself.
+ Other models had simular resolution. such that exid perturbations were roughly equivalent to velocity perturbations of 0.5% and velocity perturbation of 2.0% had a noticeable effect on the initial scale of the instability.," Other models had similar resolution, such that grid perturbations were roughly equivalent to velocity perturbations of $0.5\%$ and velocity perturbation of $2.0\%$ had a noticeable effect on the initial scale of the instability."
+ Perturbations arising from the erid were no larger than 2%. well within the regine of perturbations expected to arise from couvection.," Perturbations arising from the grid were no larger than $2\%$, well within the regime of perturbations expected to arise from convection."
+ Because the zD-series models were more compact than sA-serics models. the reverse shock reached the centers of zD inodels faster than in the sA models.," Because the zD-series models were more compact than sA-series models, the reverse shock reached the centers of zD models faster than in the sA models."
+ This had the effect of shutting off mixing in the zD-series before the Ravleigh-Tavlor instability had time to become fully nou-lnear (seealso2).., This had the effect of shutting off mixing in the zD-series before the Rayleigh-Taylor instability had time to become fully non-linear \citep[see also][]{Herant&Woosley:1994}.
+ The initial perturbations had a ercater effect on the final state of the zD simulations. while the initial perturbation scale and spectrum are washed out in the sA aodels as a result of their longer mixing times.," The initial perturbations had a greater effect on the final state of the zD simulations, while the initial perturbation scale and spectrum are washed out in the sA models as a result of their longer mixing times."
+ A simulation with random velocity perturbations of 5% was performed for Model z25D. Models were initially mapped onto the ποdimensional erid such that their inner dou cores were resolved with at least 1 blocks of 16 zones each., A simulation with random velocity perturbations of $5\%$ was performed for Model z25D. Models were initially mapped onto the two-dimensional grid such that their inner iron cores were resolved with at least 4 blocks of 16 zones each.
+ This was sufficieut to ensure that the rest of the star was accuratelv resolved., This was sufficient to ensure that the rest of the star was accurately resolved.
+ As the simulation progressed aud the moclel stars expauded. the uiximuimiu refinement leve of the simulation was turned down. so that the mode star was always resolved at about the same percentage of the radius of its iuuer core of aud heavier isotopes.," As the simulation progressed and the model stars expanded, the maximum refinement level of the simulation was turned down, so that the model star was always resolved at about the same percentage of the radius of its inner core of and heavier isotopes."
+ This is the region where Ravileigh-Tavlor müxiug takes place., This is the region where Rayleigh-Taylor mixing takes place.
+ The star expauds hoimologouslv. cusuring that al regions will be adequately resolved.," The star expands homologously, ensuring that all regions will be adequately resolved."
+" The solar composition models were mapped outo a eric 5«10"" απ on a side.", The solar composition models were mapped onto a grid $5\times 10^{14}$ cm on a side.
+ The zeroanetallicity stars were mapped onto a erid 1.1«1035 Gun ou a side., The zero-metallicity stars were mapped onto a grid $1.4\times 10^{14}$ cm on a side.
+ The portion of the erid outside the original model was initializes with a density proportional to r?., The portion of the grid outside the original model was initialized with a density proportional to $r^{-2}$.
+ Siuulatious with au outer density proportional to rH3 was also perform for Modols «ΙΑ. s25A. and z25D. These density profiles span the realistic range of inooth density distributions outside real stars.," Simulations with an outer density proportional to $r^{-3.1}$ was also performed for Models s15A, s25A, and z25D. These density profiles span the realistic range of smooth density distributions outside real stars."
+ No difference in the final profiles for density. temperature. pressure. or composition was fou yetwoeen models with differeut outer density profiles.," No difference in the final profiles for density, temperature, pressure, or composition was found between models with different outer density profiles."
+ The density profile of the surrounding material therefor has 10 effect on the amount of Ravieigh-Tavlor mixing that eoes on inside the star. provided verv little mass as a xoportion of the original mass of the star is swept up in the first davs of the explosion.," The density profile of the surrounding material therefor has no effect on the amount of Rayleigh-Taylor mixing that goes on inside the star, provided very little mass as a proportion of the original mass of the star is swept up in the first days of the explosion."
+ The amount of mass added to the erid from ambicut density was <2% for all uodels., The amount of mass added to the grid from ambient density was $< 2\%$ for all models.
+ Caleulatious were run at least until the Ravleigh-Tavlor fuegers had ceased to move with respect to the nass coordinate of the star., Calculations were run at least until the Rayleigh-Taylor fingers had ceased to move with respect to the mass coordinate of the star.
+ This happened 2 hours after core bounce for Model zi5D. 1 hours after core voce for z25D. and 7 davs after core bounce for he sl5A and s25A iocels.," This happened 2 hours after core bounce for Model z15D, 4 hours after core bounce for z25D, and 7 days after core bounce for the s15A and s25A models."
+ All models were followed o 10° seconds. long after Ravleigh-Tavlor mixing had rozen out in the zero metallicity stars. but long enough hat infall though. the immer boundary had reached an asvinptotic stage aud the fal remuaut mass from these wo-diensional simulations could be determined.," All models were followed to $10^6$ seconds, long after Rayleigh-Taylor mixing had frozen out in the zero metallicity stars, but long enough that infall though the inner boundary had reached an asymptotic stage and the final remnant mass from these two-dimensional simulations could be determined."
+ Figure 2 shows the evolution of the stable and uustable regious of the models. aud the position im mass coordinate of the forward aud reverse shocks.," Figure \ref{stability} shows the evolution of the stable and unstable regions of the models, and the position in mass coordinate of the forward and reverse shocks."
+ A reverse shock forms when the outgoiug shock encounters a region of mereasiug pi? (?7)..," A reverse shock forms when the outgoing shock encounters a region of increasing $\rho r^3$ \citep{Herant&Woosley:1994,Woosley&Weaver:1995}."
+" When the shock encounters a density profile that falls off with a flatter slope than xr3 Ίο, when it cucomuters a reeion of increasing pi? it decelerates."," When the shock encounters a density profile that falls off with a flatter slope than $\propto r^{-3}$, i.e. when it encounters a region of increasing $\rho r^3$ it decelerates."
+ The deceleration of the forward shock front reverses the direction of the pressure eracicut. which slows down the lavers interior to the shock. as well.," The deceleration of the forward shock front reverses the direction of the pressure gradient, which slows down the layers interior to the shock, as well."
+ Shocked material piles up and forms a lieh deusity post-shock shell., Shocked material piles up and forms a high density post-shock shell.
+ The reverse shock forms at the inner boundary of the high-density shell of decelerated matter and propagates down into the star. toward its center. slowing down the deeper. inner lavers of the star (?)..," The reverse shock forms at the inner boundary of the high-density shell of decelerated matter and propagates down into the star, toward its center, slowing down the deeper, inner layers of the star \citep{Kifonidis:2003}."
+ The deceleration of the shock creates a steep pressure eracicut in the opposite direction to the existing eravitational aud density ΠΕ, The deceleration of the shock creates a steep pressure gradient in the opposite direction to the existing gravitational and density gradients.
+ This pressure eracdient can overwhelui he gravitational eradient. and in doiug so trigecrs he formation of Ravleigh-Tavlor iustabilities in the uaterial.," This pressure gradient can overwhelm the gravitational gradient, and in doing so triggers the formation of Rayleigh-Taylor instabilities in the material."
+ The Ravleigh-Tavlor instability develops uutil he reverse shock has passed by. at which poiut the uaterial becomes stable again. aud the instabilitics cease o grow exponeutial.," The Rayleigh-Taylor instability develops until the reverse shock has passed by, at which point the material becomes stable again, and the instabilities cease to grow exponentially."
+ Figure 2. covers the period of ine from when the models were first mapped to o sheltly bevoud the time when the Ravleigh-Tavlor," Figure \ref{stability}
+ covers the period of time from when the models were first mapped to to slightly beyond the time when the Rayleigh-Taylor"
+ol the CDF comparisons are listed in Table 2.,of the CDF comparisons are listed in Table 2.
+ As in the traditional IXolomogorov-9mirnov test. small values of the sienilicance indicate that the CDFs are incompatible.," As in the traditional Kolomogorov-Smirnov test, small values of the significance indicate that the CDFs are incompatible."
+ From these significances. we infer the following: (1) The SAD areas are consistent only with a log-normal distribution. (," From these significances, we infer the following: (1) The SAD areas are consistent only with a log-normal distribution. ("
+2) The SAD fluxes are inconsistent with a normal or power-law distribution. but consistent with both log-normale and exponential distributions.,"2) The SAD fluxes are inconsistent with a normal or power-law distribution, but consistent with both log-normal and exponential distributions."
+ As an additional test. we repeated the area analvsis omitting the SADs measured in the 26 Jun 2001 flare.," As an additional test, we repeated the area analysis omitting the SADs measured in the 26 Jun 2001 flare."
+ This flare contributes the majoritv of SAD areas smaller than 20 Mnpr. and therefore one may reasonably ask if the overall distribution is affected by their inclusion/exchision.," This flare contributes the majority of SAD areas smaller than 20 $^2$, and therefore one may reasonably ask if the overall distribution is affected by their inclusion/exclusion."
+ We find that the result is unchanged: the areas of the other 111 SADs are consistent wilh a log-normal distribution (with Kuiper significance 0.4). ancl (he other distributions are ruled out completely (significances between 5x10. and 6x10 3).," We find that the result is unchanged: the areas of the other 111 SADs are consistent with a log-normal distribution (with Kuiper significance 0.4), and the other distributions are ruled out completely (significances between $5\times10^{-4}$ and $6\times10^{-31}$ )."
+ Finally. we examined the 25 SADs from the 20 Jan 1999 flare (i.e... Group 3).," Finally, we examined the 25 SADs from the 20 Jan 1999 flare (i.e., Group 3)."
+ Though the sanmiple is much smaller. we find (hat (he areas are consistent wilh a loe-normal distribution (significance 0.3). and less likely to be normally distributed (signif.," Though the sample is much smaller, we find that the areas are consistent with a log-normal distribution (significance 0.3), and less likely to be normally distributed (signif."
+ 8x10. 2)., $8\times10^{-2}$ ).
+ Power-law and exponential distributions are discounted with signilicances 2x10. and 6xI0.. respectively.," Power-law and exponential distributions are discounted with significances $2\times10^{-3}$ and $6\times10^{-8}$, respectively."
+ Although the TRACE data of Savage&Melxenzie(2011) provide size and fIux estimates for some 23 SADs. the cumulative distribution curves of the TRACE data alone are very noisy. due to the sparseness of the sample.," Although the TRACE data of \citet{SavageMcKenzie_11} provide size and flux estimates for some 23 SADs, the cumulative distribution curves of the TRACE data alone are very noisy, due to the sparseness of the sample."
+ The TRACE measurements indicate that SADs smaller (han SXTs resolution exist. so one can conclude that the Irequency. distribution below the range shown in Figure 2 is non-zero.," The TRACE measurements indicate that SADs smaller than SXT's resolution exist, so one can conclude that the frequency distribution below the range shown in Figure 2 is non-zero."
+ However. combination of the SXT and TRACE," However, combination of the SXT and TRACE"
+of the characteristics of the (incomplete) large population of bubbles reported.,of the characteristics of the (incomplete) large population of bubbles reported.
+ A later catalog for bubbles within |/|<10° in the Galaxy. found à comparable number., A later catalog for bubbles within $|l| < 10^\circ$ in the Galaxy found a comparable number.
+ These bubbles were identified by rings of 8j/m emission that were found to enclose 24jm emission; a few such infrared bubbles had been detected before by the and by the2??)., These bubbles were identified by rings of $8\um$ emission that were found to enclose $24\um$ emission; a few such infrared bubbles had been detected before by the and by the.
+. N49 was later the focus of a more detailed study by?., N49 was later the focus of a more detailed study by.
+. For that study. N49 was chosen because of its relatively symmetric appearance: for instance. the 8j/m emission was very well fit by assuming a spherical shell?).," For that study, N49 was chosen because of its relatively symmetric appearance: for instance, the $8\um$ emission was very well fit by assuming a spherical shell."
+. This implies that perhaps N49 is à useful test case for comparing with models of WBB evolution., This implies that perhaps N49 is a useful test case for comparing with models of WBB evolution.
+ also examined 24jim emission observed with the Multiband Imaging Photometer aboard as part of the the MIPS Galactic Plane Survey ?).. comparing that data to 8m observations and to cem radio observations(?).," also examined $24\um$ emission observed with the Multiband Imaging Photometer aboard as part of the the MIPS Galactic Plane Survey , comparing that data to $8\um$ observations and to cm radio observations."
+. These data showed that emission at 24j/m and 20cem are coincident. sharing the same annular geometry (although the 20 em data. indicating free-free emission from ionized gas. does not drop off quite as quickly with radius as the 24j/m emission): both were also enclosed by the larger-radius 8jm emission of the photodissociation region (PDR) shell. strongly indicating that the dusty post-shocked gas resides within the outlying 8m PDR shell.," These data showed that emission at $24\um$ and cm are coincident, sharing the same annular geometry (although the 20 cm data, indicating free-free emission from ionized gas, does not drop off quite as quickly with radius as the $24\um$ emission); both were also enclosed by the larger-radius $8\um$ emission of the photodissociation region (PDR) shell, strongly indicating that the dusty post-shocked gas resides within the outlying $8\um$ PDR shell."
+ However. it is not at all clear if a reasonable mass of dust could explain the observed emission. and if dust could survive in that post-shocked gas environment.," However, it is not at all clear if a reasonable mass of dust could explain the observed emission, and if dust could survive in that post-shocked gas environment."
+ And. if dust can exist within the bubble. what role might such dust play in the evolution of wind-blown bubbles?," And, if dust can exist within the bubble, what role might such dust play in the evolution of wind-blown bubbles?"
+" These are the questions that the present paper sets out to answer,", These are the questions that the present paper sets out to answer.
+ Our first task was to see whether dust emission (alone) is a reasonable fit to the observations., Our first task was to see whether dust emission (alone) is a reasonable fit to the observations.
+" To synthesize observations of dust in WBBs. we have run phototonization simulations using ""Cloudy""?)."," To synthesize observations of dust in WBBs, we have run photoionization simulations using `Cloudy'."
+. In those Cloudy simulations. we investigated emission from a constant-temperature ISM plasma with an included “normal” dust grain population with a dust-mass-to-gas-mass ratio of 6.4«107.," In those Cloudy simulations, we investigated emission from a constant-temperature ISM plasma with an included “normal” dust grain population with a dust-mass-to-gas-mass ratio of $6.4\times10^{-3}$."
+ We employed Cloudy_33D to map the ID Cloudy simulations into a spherically-symmetric— 3D bubble that was then projected onto the plane of the sky to compare with observations., We employed 3D to map the 1D Cloudy simulations into a spherically-symmetric 3D bubble that was then projected onto the plane of the sky to compare with observations.
+ Note that Cloudy cannot model the dynamic post-shock structure of the WBB. so for photoionization studies. we model the impact of dusty gas on a bubble of uniform temperature which we set.," Note that Cloudy cannot model the dynamic post-shock structure of the WBB, so for photoionization studies, we model the impact of dusty gas on a bubble of uniform temperature which we set."
+" At this early stage. we have no dynamical model of dust fluxes into. or out of. the bubble: we are examining the emission from a dusty ISM plasma within the bubble. with a standard ISM dust-to-gas ratio,"," At this early stage, we have no dynamical model of dust fluxes into, or out of, the bubble; we are examining the emission from a dusty ISM plasma within the bubble, with a standard ISM dust-to-gas ratio."
+ The parameters for the simulations and analysis are summarized in Table 1.. along with references for the values used.," The parameters for the simulations and analysis are summarized in Table \ref{cloudyModelParams}, along with references for the values used."
+ We describe and justify those parameter choices in the paragraphs that follow., We describe and justify those parameter choices in the paragraphs that follow.
+" For the central star and its central stellar continuum. we include a star of class OSV(2).. Using Cloudy’s ""wmbasic? stellar continuum modelswinds).. and calling on the models in Table 1 of?.. we set the star’s luminosity to 107?!ZL... the surface temperature to Tay=41540 K. the surface gravity to log(g)=3.92. and the star's metallicity to solar."," For the central star and its central stellar continuum, we include a star of class O5V. Using Cloudy's `wmbasic' stellar continuum models, and calling on the models in Table 1 of, we set the star's luminosity to $10^{5.51}~L_{\odot}$, the surface temperature to $T_{\rm eff} =
+41540$ K, the surface gravity to $\log(g) = 3.92$, and the star's metallicity to solar."
+ These parameters correspond to a 37.3M.. star. with a flux of hydrogen-1onizing photons of logqo=107? photons em™ s or logQa=107° photons s7' into 47 steradians. very close to the value cited by?.," These parameters correspond to a $37.3~M_{\odot}$ star, with a flux of hydrogen-ionizing photons of $\log q_0 = 10^{24.4}$ photons $^{-2}$ $^{-1}$ or $\log Q_0 =
+10^{49.3}$ photons $^{-1}$ into $4\pi$ steradians, very close to the value cited by."
+. The stellar mass loss rate. Maing is not used directly in the Cloudy simulations. but used in our analysis of the velocity of the post-shocked gas to estimate (later) the dust-gas friction that drags dust out of the post-shocked wind region.," The stellar mass loss rate, $\dot{M}_{\rm wind}$ is not used directly in the Cloudy simulations, but used in our analysis of the velocity of the post-shocked gas to estimate (later) the dust-gas friction that drags dust out of the post-shocked wind region."
+" For the stellar parameters we've adopted. the IDL code ""Mdot.pro' from yields a mass outflow rate of 1.3«1079 .. vi! and va~2600 km s."," For the stellar parameters we've adopted, the IDL code `Mdot.pro' from yields a mass outflow rate of $1.4\times10^{-6}$ $_{\odot}$ $^{-1}$ and $v_{\infty} \sim
+2600$ km $^{-1}$."
+ This yields a wind kinetic luminosity of 3<10° ergs sl.," This yields a wind kinetic luminosity of $3 \times
+10^{36}$ ergs $^{-1}$."
+ The kinematic distance to N49 is 5.73:0.6 kpe(2)., The kinematic distance to N49 is $5.7 \pm 0.6$ kpc.
+. Given this distance. the size of the bubble. roue;. is set by the inner radius of PAH emission. so ts observationally estimated as ppe.," Given this distance, the size of the bubble, $r_{\rm outer}$, is set by the inner radius of PAH emission, so is observationally estimated as pc."
+ The density of ambient gas surrounding N49 is not well constrained. but unpublished observations of NH; emission (Cyganowski et al..," The density of ambient gas surrounding N49 is not well constrained, but unpublished observations of ${\rm NH}_3$ emission (Cyganowski et al.,"
+ in preparation) surrounding N49 indicate a high ambient gas density of order 10 cem™., in preparation) surrounding N49 indicate a high ambient gas density of order $10^4$ $^{-3}$ .
+ The age of N49 is not well constrained. but if the external medium has a gas density of order 10 em™. the age of a wind-blown bubble of the inferred size for N49 is approximately 5«10? to 10° years(2).," The age of N49 is not well constrained, but if the external medium has a gas density of order $10^4$ $^{-3}$, the age of a wind-blown bubble of the inferred size for N49 is approximately $5 \times 10^5$ to $10^6$ years."
+. Again. since Cloudy does not simulate the heating of the wind's post-shocked gas. we start by modelling the WBB with a uniform temperature for the post-shocked gas.," Again, since Cloudy does not simulate the heating of the wind's post-shocked gas, we start by modelling the WBB with a uniform temperature for the post-shocked gas."
+ For most of the simulations. we have chosen 3.5«10° K as that constant temperature: this value for the temperature is informed by our own simple dynamical models of wind-blown bubble dynamies with dust cooling (see Section ??)).," For most of the simulations, we have chosen $3.5\times10^6$ K as that constant temperature; this value for the temperature is informed by our own simple dynamical models of wind-blown bubble dynamics with dust cooling (see Section \ref{modelResults}) )."
+ Changes in the WBB temperature affect the amount of dust that is inferred: increasing the temperature to 10’ KK as suggested in the dust-free models of would reduce the amount of dust required. but not eliminate it.," Changes in the WBB temperature affect the amount of dust that is inferred; increasing the temperature to $10^7$ K as suggested in the dust-free models of would reduce the amount of dust required, but not eliminate it."
+ Finally. we assume that the grain distribution. which we specify inCloudy. ts an ISM-type grain distribution (??)..We explain below. in Section ?? why that grain distribution seems to be preferred by the observations.," Finally, we assume that the grain distribution, which we specify inCloudy, is an ISM-type grain distribution .We explain below, in Section \ref{cloudyModelFits}
+ why that grain distribution seems to be preferred by the observations."
+ First. we ask whether dust emission canreproduce the 24ym observations of N49.," First, we ask whether dust emission canreproduce the $\um$ observations of N49."
+ Using the above-defined, Using the above-defined
+Fluctuation operator and mode functions,The potential ${\cal V}$ will be specified below after partial wave decomposition.
+ Thefluctuation operatoris definedingeneral field wherec; denoteshere thethesefluctuating willbefieldsthe and ot!theand“classical” background For our specific modelwe expandas dMοἱ (3.1). Ay ca... i Ley = =51D (Outta)42 −4 qDcl cls «n (cls," The fluctuation operator $\calm$ can be decomposed into partial waves and its determinant decomposes accordingly, We introduce the following partial wave decomposition for fields After inserting these expressions into the Lagrange density and using the reality conditions for the fields one finds that the following combinations are real relative to each other and make the fluctuation operators symmetric: Writing the partial fluctuation operators - omitting the index $n$ in the following - as the free operators ${\bf M}^0$ become diagonal matrices with elements"
+simular to the radio core in III Zw 2 which was the first Sevfert ealaxy discovered to contain a superluminal jet.,similar to the radio core in III Zw 2 which was the first Seyfert galaxy discovered to contain a superluminal jet.
+ This galaxy has a müillimieter-peaked. spectimu and a jet which shows a stop-aud-eo behavior indicative of a strong interaction with deuse material on the sub-parsec scale (Falckeetal.1999:Bruuthaler 20000).," This galaxy has a millimeter-peaked spectrum and a jet which shows a stop-and-go behavior indicative of a strong interaction with dense material on the sub-parsec scale \cite{FalckeBowerLobanov1999,BrunthalerFalckeBower2000}. ."
+ Ulvestadetal.(1999) therefore speculate whether the bright inverted radio core in Ak 2318 could be interpreted simularlv to hose iu Compact Syxauunetrie Objects (CSOs) with a Cüsgalertz-Peuked-Spectiuu (GPS. see O'Dea1998)).," \citeN{UlvestadWrobelRoy1999}
+ therefore speculate whether the bright inverted radio core in Mrk 348 could be interpreted similarly to those in Compact Symmetric Objects (CSOs) with a Gigahertz-Peaked-Spectrum (GPS, see \citeNP{O'Dea1998}) )."
+ Iu hese galaxies bright hotspots are formed im a jet that cruniuates already on the parsec scale., In these galaxies bright hotspots are formed in a jet that terminates already on the parsec scale.
+ In HII Zw 2 aud Ak 318 this secs o happen on even smaller scales. cading to higher peak frequencies and could be due to rustration of the jet by a molecular cloud or even a warped or mnisaliened torus.," In III Zw 2 and Mrk 348 this seems to happen on even smaller scales, leading to higher peak frequencies and could be due to frustration of the jet by a molecular cloud or even a warped or misaligned torus."
+ Since the masers in NGC 1052. which have similar woad line widths as in Mrk 3148. are founcd aloug the radio jet (Claussenctal.1905) it should be checked whether in Mrk 318 one has an analogous situation.," Since the masers in NGC 1052, which have similar broad line widths as in Mrk 348, are found along the radio jet \cite{ClaussenDiamondBraatz1998} it should be checked whether in Mrk 348 one has an analogous situation."
+ One can speculate that iu such a case the evolution of the radio flare and the evolution of the mascr flare aud its blue wing could be related. possibly providing a unique diagnostic tool to study jet-ISA interactions.," One can speculate that in such a case the evolution of the radio flare and the evolution of the maser flare and its blue wing could be related, possibly providing a unique diagnostic tool to study jet-ISM interactions."
+ Tn any case. with its bright radio core Mrk 318 provides an ideal opportunity to observe the maser lines m this ealaxy at high resolution with VLBI during this flare even though the lines still have a rather low flux.," In any case, with its bright radio core Mrk 348 provides an ideal opportunity to observe the maser lines in this galaxy at high resolution with VLBI during this flare even though the lines still have a rather low flux."
+ Since radio and maser emission seen to be highly variable both should be monitored frequeuth., Since radio and maser emission seem to be highly variable both should be monitored frequently.
+ Caven that Ak 2318 was not discovered in an earlier survey this finding also suggests that existing samples should be revisited to search for more flaringnieganuasers., Given that Mrk 348 was not discovered in an earlier survey this finding also suggests that existing samples should be revisited to search for more flaringmegamasers.
+flux ratios and we never use them to discuss the difference in the properties of the NLR emissiou between NLSIs and BLSIs.,flux ratios and we never use them to discuss the difference in the properties of the NLR emission between NLS1s and BLS1s.
+ The gas inctallicity in NLRs of $2s has been sometimes studied by using the ciiissiou-line fux ratio of [N UJAG583/Tla (e.g.. Storchi-Beremaun Pastoriza 1989. 1990: Storcli-Beremanu 1991: Radovich Rafauclli 1996).," The gas metallicity in NLRs of S2s has been sometimes studied by using the emission-line flux ratio of [N $\lambda$ $\alpha$ (e.g., Storchi-Bergmann Pastoriza 1989, 1990; Storchi-Bergmann 1991; Radovich Rafanelli 1996)."
+ Using photoionization models. Storchi-Dergmaun Pastoriza (1989. 1990) showed that this emission-line flux ratio is sensitive to the nitroseu abuudanuce but rather insensitive to the eas deusitv and the ionization parameter (G.c.. a uunboer density ratio of jonizing photous to hydrogen atoms).," Using photoionization models, Storchi-Bergmann Pastoriza (1989, 1990) showed that this emission-line flux ratio is sensitive to the nitrogen abundance but rather insensitive to the gas density and the ionization parameter (i.e., a number density ratio of ionizing photons to hydrogen atoms)."
+ Note that the nitrogen abundance is portant since the N/O abuudauce ratio in galactic ITnu regions is known to scale with the O/T abundance ratio (c.e@.. Shields 1976: Pagel οι 1981: Villa-Costas Ediimnuds 1993: van Zee. Salzer Tavucs 1998: Izotov Thuan 1999).," Note that the nitrogen abundance is important since the N/O abundance ratio in galactic H regions is known to scale with the O/H abundance ratio (e.g., Shields 1976; Pagel Edmunds 1981; Villa-Costas Edmunds 1993; van Zee, Salzer Haynes 1998; Izotov Thuan 1999)."
+ Therefore the nitrogen abundance. N/TL scales with Z?. where Z is the gas iietallicitvy.," Therefore the nitrogen abundance, N/H, scales with $Z^2$, where $Z$ is the gas metallicity."
+ The data of IT regions show that these scaling relations hold approximately for Z>0.22..., The data of H regions show that these scaling relations hold approximately for $Z \gtrsim 0.2 Z_{\odot}$.
+ We show the histograms of the cuiissiou-line flux ratio of [N WJAG5S3/Tonanos for the NLS1s aud the DESIs in the sample of Rodríeeuez-Ardila et al. (, We show the histograms of the emission-line flux ratio of [N $\lambda$ $\alpha_{\rm narrow}$ for the NLS1s and the BLS1s in the sample of guez-Ardila et al. (
+2000) in Figure lL.,2000) in Figure 1.
+ The data are not corrected for the dust extinction because the effect of the dust extinction. ou this flux ratio is neslieiblv sinall., The data are not corrected for the dust extinction because the effect of the dust extinction on this flux ratio is negligibly small.
+ The average aud median values of [N nA6583/Ilouae for the NLS1s aud those for the DLS1s are eiven in Table 1., The average and median values of [N $\lambda$ $\alpha_{\rm narrow}$ for the NLS1s and those for the BLS1s are given in Table 1.
+ There is no appareut difference in this flux ratio between the two populations., There is no apparent difference in this flux ratio between the two populations.
+ It is also suggested bv the EFKolnogorov-Suirnov (199) statistical test that the two frequency distributious of the flux ratio of [N HJAG5S3/Tlenance are statistically indistinenishable (Py= 0.912)., It is also suggested by the Kolmogorov-Smirnov (KS) statistical test that the two frequency distributions of the flux ratio of [N $\lambda$ $\alpha_{\rm narrow}$ are statistically indistinguishable $P_{\rm KS} = 0.912$ ).
+ As shown in Figure 2. we cannot find an apparent correlation between the fux ratio of [N Πλ ΠΠgeno and full-width at half παΤΗ (FWIIM) of the broad component of the IIo. emission.," As shown in Figure 2, we cannot find an apparent correlation between the flux ratio of [N $\lambda$ $\alpha_{\rm narrow}$ and full-width at half maximum (FWHM) of the broad component of the $\alpha$ emission."
+ The correspouding correlation cocficicut is 0.009. which iuplies that there is no meaningful correlation between the two quautitics.," The corresponding correlation coefficient is 0.009, which implies that there is no meaningful correlation between the two quantities."
+ This result secms contrary to the result of Wills et al (, This result seems contrary to the result of Wills et al. (
+1999). who fouud an appareut uegative correlation between the cuussiou-line flux ratio of N A1210/C ni]A1909. which is sensitive to the nitrogen abundance. aud the PFWIIM. of the broad component of the IJ cussion.,"1999), who found an apparent negative correlation between the emission-line flux ratio of N $\lambda$ 1240/C $\lambda$ 1909, which is sensitive to the nitrogen abundance, and the FWHM of the broad component of the $\beta$ emission."
+ Based oon the intensive photoionization model calculations and the careful calibrations with observations. Storchi-Beremann et al. (," Based on the intensive photoionization model calculations and the careful calibrations with observations, Storchi-Bergmann et al. ("
+1998. hereafter SSCI) proposed the revised diagnostics for the gas metallicity of NLRs in ACNs.,"1998, hereafter SSCK) proposed the revised diagnostics for the gas metallicity of NLRs in AGNs."
+ They showed that the oxveen abundance of the NLR eas in ACNs can be expressed by. where a— [N ajA65s3/llouuusse and yo= 1ο ABOUT asso. and where à= log (JO n|A3727/|O. A5007) and e= log QN n[AG583/TIo saos).," They showed that the oxygen abundance of the NLR gas in AGNs can be expressed by where $x \equiv$ [N $\lambda$ $\alpha_{\rm narrow}$ and $y \equiv$ [O $\lambda$ $\beta_{\rm narrow}$, and where $u \equiv$ log ([O $\lambda$ 3727/[O $\lambda$ 5007) and $v \equiv$ log ([N $\lambda$ $\alpha_{\rm narrow}$ )."
+ These methods weakly depend on the eas density: the derived values should be subtracted, These methods weakly depend on the gas density; the derived values should be subtracted
+Molecular clouds are lighly filamentary structures that are supported against rapid. global collapse by a combination of ordered maeuctic fields as well as by uou-hermal. hydromagnuetie turbulence of some kind.,"Molecular clouds are highly filamentary structures that are supported against rapid, global collapse by a combination of ordered magnetic fields as well as by non-thermal, hydromagnetic turbulence of some kind."
+ However. determining the structure of the field has proven difficult.," However, determining the structure of the field has proven difficult."
+ Direct Zeeiiun mcasurcuicuts have been limited to a few points per cloud rather than full-scaο luaps. uid rarely trace the highest density eas (Crutcher et al.," Direct Zeeman measurements have been limited to a few points per cloud rather than full-scale maps, and rarely trace the highest density gas (Crutcher et al."
+ 1993. 1996. 1999).," 1993, 1996, 1999)."
+ Polarization maps in οofieal aud near-infrared extinction fail to reveal the structure of the field in the deuse gas. perhaps as a result ofthe poor polarizing power of grains at these waveleneths in regions of high optical depth (Goodima vet al.," Polarization maps in optical and near-infrared extinction fail to reveal the structure of the field in the dense gas, perhaps as a result of the poor polarizing power of grains at these wavelengths in regions of high optical depth (Goodman et al."
+ 1995)., 1995).
+ Observations at far-infrared aud sub-millimetre waveleueths have demoustrated that the thermal Cluission from dust eras is often partially polarized (ee., Observations at far-infrared and sub-millimetre wavelengths have demonstrated that the thermal emission from dust grains is often partially polarized (eg.
+ Scheuniue 1998. Matthews Wilson 2000.," Schleuning 1998, Matthews Wilson 2000,"
+On the basis of the above. one may conclude that rere is no uniformity in extinction properties amongst these usters.,"On the basis of the above, one may conclude that there is no uniformity in extinction properties amongst these clusters."
+ In fact. they cüller from cluster to cluster.," In fact, they differ from cluster to cluster."
+ ALL this indicates that non-uniform extinction. observed in. voung clusters cannot be understood in terms of a simple physical scenario., All this indicates that non-uniform extinction observed in young clusters cannot be understood in terms of a simple physical scenario.
+ Actually. it depends upon factors such as the age of the cluster members. initial spatial distribution of matter in the molecular clouds. sequential star formation processes and the distribution of hot O and D stars in the cluster etc.," Actually, it depends upon factors such as the age of the cluster members, initial spatial distribution of matter in the molecular clouds, sequential star formation processes and the distribution of hot O and B stars in the cluster etc."
+ A complicated physical scenario is therefore required to explain the non-uniform extinction in voung star clusters., A complicated physical scenario is therefore required to explain the non-uniform extinction in young star clusters.
+ We are grateful to. the referee. Dr. Alike Bessell. for valuable comments which improved. the scientifie content of the paper.," We are grateful to the referee, Dr. Mike Bessell, for valuable comments which improved the scientific content of the paper."
+ Useful discussions with Dr. A. Ix. Pandev. is thankfully acknowledged., Useful discussions with Dr. A. K. Pandey is thankfully acknowledged.
+with a standard deviation of σ~0.015 dex around this relation.,with a standard deviation of $\sigma \sim 0.015$ dex around this relation.
+" We use the rest-frame g—r colour as a proxy for stellar continuum, but other colours such as r—i provide similar constraints."," We use the rest-frame $g-r$ colour as a proxy for stellar continuum, but other colours such as $r-i$ provide similar constraints."
+" In the case of low-S/N spectra, the rest-frame colour could be from SED fitting to broad-band magnitudes."," In the case of low-S/N spectra, the rest-frame colour could be from SED fitting to broad-band magnitudes."
+" Thus, in summary, the Balmer decrement for a low-resolution, strong-emission-line spectrum can be measured from the emission line equivalent widths and colours alone with; with a scatter around this of σ~0.05 dex, or ~0.3 mag in Ay, assuming a Galactic extinction law (e.g.O'Donnell 1994)."," Thus, in summary, the Balmer decrement for a low-resolution, strong-emission-line spectrum can be measured from the emission line equivalent widths and colours alone with; with a scatter around this of $\sigma\sim0.05$ dex, or $\sim0.3$ mag in $A_{V}$ , assuming a Galactic extinction law \citep[e.g.][]{ODonnell94}."
+". For weaker emission line galaxies ((i.e. EW(Ho))«30A)), a lower offset (R~ 3.5) should be used, as shown in figure 4.."," For weaker emission line galaxies ((i.e. $<30$ ), a lower offset $R\sim3.5$ ) should be used, as shown in figure \ref{fig:binEWfix}."
+" However, given the larger uncertainties and greater dispersion seen at lower EW(Ha)), a correction factor of R~4 can be used for the full sample with a 0.1 dex scatter (~0.7 mag in Ay) and an extension to lower values (i.e EW Balmer decrement underestimate) due to low EW systems."," However, given the larger uncertainties and greater dispersion seen at lower ), a correction factor of $R\sim4$ can be used for the full sample with a 0.1 dex scatter $\sim0.7$ mag in $A_{V}$ ) and an extension to lower values (i.e EW Balmer decrement underestimate) due to low EW systems."
+" One final note on this relation: as seen in figure 3,, there is a strong bias in the sample of Balmer decrement with other galaxy properties, as discussed in detail in several other SDSS papers (seee.g.Kauffmannetal.2003a;Garn&Best 2010)."," One final note on this relation: as seen in figure \ref{fig:EW_HaHb}, there is a strong bias in the sample of Balmer decrement with other galaxy properties, as discussed in detail in several other SDSS papers \citep[see
+e.g.][]{Kauffmann03a,Garn10}."
+". Thus, the relation shown above includes a combination of both galaxy type as well as variation in extinction."," Thus, the relation shown above includes a combination of both galaxy type as well as variation in extinction."
+" The only way to remove fully this effect is to match pairs of galaxies in as many property types excluding extinction, such as done in Wildetal.(2011b)."," The only way to remove fully this effect is to match pairs of galaxies in as many property types excluding extinction, such as done in \citet{Wild11b}."
+". Unfortunately when applied to the sample here, it was found that the range in extinction was not large enough to properly determine the relation."," Unfortunately when applied to the sample here, it was found that the range in extinction was not large enough to properly determine the relation."
+" However, even given these uncertainties, this relation should still hold at several redshifts as high attenuations are on average associated with high gas masses, and thus high star formation rates and similar underlying continua at all redshifts."," However, even given these uncertainties, this relation should still hold at several redshifts as high attenuations are on average associated with high gas masses, and thus high star formation rates and similar underlying continua at all redshifts."
+ One issue with the previous section is that we assume throughout that the stellar-continuum subtracted emission line fluxes within the MPA/JHU database are correct., One issue with the previous section is that we assume throughout that the stellar-continuum subtracted emission line fluxes within the MPA/JHU database are correct.
+" While the overall fits to the stellar continuum are impressively good with a median y?=1.01 per pixel across the sample, there are appear to be remaining issues around the Balmer lines."," While the overall fits to the stellar continuum are impressively good with a median $\chi^2 = 1.01$ per pixel across the sample, there are appear to be remaining issues around the Balmer lines."
+" In the following we concentrate on the SDSS DR7 Balmer emission-line fluxes corrected for the underlying stellar absorption features from the MPA/JHU database, and explore their uncertainties using the known intrinsic values and commonly used attenuation and extinction laws."," In the following we concentrate on the SDSS DR7 Balmer emission-line fluxes corrected for the underlying stellar absorption features from the MPA/JHU database, and explore their uncertainties using the known intrinsic values and commonly used attenuation and extinction laws."
+" When considered alone, the ratio of ccannot indicate problems with the measurement of the lines involved unless it is significantly below the expected value of the unattenuated ratio."," When considered alone, the ratio of cannot indicate problems with the measurement of the lines involved unless it is significantly below the expected value of the unattenuated ratio."
+" This is because the larger values of the emission line ratio can be caused by attenuation by intervening dust, with the intrinsic ratio dependent the emitting gas density and temperature (as discussed in section ??))."," This is because the larger values of the emission line ratio can be caused by attenuation by intervening dust, with the intrinsic ratio dependent the emitting gas density and temperature (as discussed in section \ref{sec:balmer}) )."
+ However by examining several of the Balmer lines at once these dependencies can be accounted for., However by examining several of the Balmer lines at once these dependencies can be accounted for.
+" Figure 6 shows the variation of the rratio against the rratio for the SN(Ho,,H8,Hy)) SDSS sample.", Figure \ref{fig:HaHbHgHb} shows the variation of the ratio against the ratio for the ) SDSS sample.
+" Both ratios have been normalized to their Case B, T210,000K, n,=100 intrinsic ratios ((Ho/H8);=2.86, (Hy/HB)= 0468)."," Both ratios have been normalized to their Case B, T=10,000K, $n_{e}=100$ intrinsic ratios $(\ha/\hb)_{\rm I}=2.86$, $(\hg/\hb)_{\rm I}=0.468$ )."
+" Immediately obvious in this figure is the offset of the sample from the zero point, indicating most SDSS galaxies undergo some attenuation (as seen in the previous plots), with the correlation between the two ratios as expected from the reddening laws applied to the intrinsic ratio."," Immediately obvious in this figure is the offset of the sample from the zero point, indicating most SDSS galaxies undergo some attenuation (as seen in the previous plots), with the correlation between the two ratios as expected from the reddening laws applied to the intrinsic ratio."
+ The three different lines overplotted show the effect of three different attenuation laws commonly assumed in the analysis of galaxies., The three different lines overplotted show the effect of three different attenuation laws commonly assumed in the analysis of galaxies.
+" For all three lines, the symbols indicates steps of 0.5 in Ay, up to Ay=3."," For all three lines, the symbols indicates steps of 0.5 in $A_{V}$, up to $A_{V} = 3$."
+" The O'Donnell(1994) law (O'D94) is an updated version of the Cardellietal.(1989) fit to the average extinction law in the Galaxy, thus least representative of the integrated emission from the SDSS galaxies, which will suffer attenuation due to the mixture of emitting sources and absorbing medium."," The \citet{ODonnell94} law (O'D94) is an updated version of the \citet{Cardelli89} fit to the average extinction law in the Galaxy, thus least representative of the integrated emission from the SDSS galaxies, which will suffer attenuation due to the mixture of emitting sources and absorbing medium."
+" However, as discussed in Kennicuttetal.(2009) and can be seen in figure 6,, the use of a foreground dust screen with galactic extinction is indistinguishable from the other laws, especially given the uncertainty within the SDSS sample."," However, as discussed in \citet{Kennicutt09} and can be seen in figure \ref{fig:HaHbHgHb}, the use of a foreground dust screen with galactic extinction is indistinguishable from the other laws, especially given the uncertainty within the SDSS sample."
+" We assume a total to selective V-band extinction of Ry=3.1, the average value in our galaxy."," We assume a total to selective $V$ -band extinction of $R_{V}=3.1$, the average value in our galaxy."
+" The Calzettietal.(2000) attenuation law (Cal00) was obtained from the continuum and Balmer decrement of local actively star-forming galaxies, thus matching the high EW(Ho)) galaxies in the sample."," The \citet{Calzetti00} attenuation law (Cal00) was obtained from the continuum and Balmer decrement of local actively star-forming galaxies, thus matching the high ) galaxies in the sample."
+" Note that as only ratios are analysed here, the difference between the colour excess (E(B— V)) of the stellar continuum and nebular lines noted by Calzetti et iis effectively scaled out."," Note that as only ratios are analysed here, the difference between the colour excess $E(B-V)$ ) of the stellar continuum and nebular lines noted by Calzetti et is effectively scaled out."
+"The Ry used here is 4.05, as given by Calzettietal. from the comparison of the observed infrared flux tothat predicted from the obscuration of the optical-ultraviolet light.","The $R_{V}$ used here is 4.05, as given by \citet{Calzetti00} from the comparison of the observed infrared flux tothat predicted from the obscuration of the optical-ultraviolet light."
+"The saturation of to D""~3-4 (Figure 16)) in analogy with that of the D2(M)Lorenz system with noise (Figure 7), implies that the temporal behavior of the system is governed by 3-4 underlying global equations.","The saturation of $D_2 (M)$ to $D_2^{sat} \approx 3-4$ (Figure \ref{GRS_beta_d2}) ) in analogy with that of the Lorenz system with noise (Figure \ref{lor_D2_noise}) ), implies that the temporal behavior of the system is governed by 3-4 underlying global equations."
+" On the other hand, it is believed that the temporal evolution of an accretion disk can be represented by local magneto-hydrodynamic equations which are non-linear differential equations in both time and disk radius."," On the other hand, it is believed that the temporal evolution of an accretion disk can be represented by local magneto-hydrodynamic equations which are non-linear differential equations in both time and disk radius."
+" Thus, an encouraging aspect of the results obtained here is that these complicated magneto-hydrodynamic equations may be reducible to (or at least approximated by) a set of simple (although still non-linear) global equations, which are functions of global parameters like size of the disk, average density, pressure etc."," Thus, an encouraging aspect of the results obtained here is that these complicated magneto-hydrodynamic equations may be reducible to (or at least approximated by) a set of simple (although still non-linear) global equations, which are functions of global parameters like size of the disk, average density, pressure etc."
+" Such a set of global equations, will provide a complete picture of the global properties of the inner accretion disk and in principle may enable the testing of general relativity in the strong field limit."," Such a set of global equations, will provide a complete picture of the global properties of the inner accretion disk and in principle may enable the testing of general relativity in the strong field limit."
+" However, obtaining these equations from the magneto-hydrodynamic equations directly is challenging."," However, obtaining these equations from the magneto-hydrodynamic equations directly is challenging."
+ Perhaps spatial averaging of these equations to obtain a one (or two zone) models like as done by Paczynski(1983) may be the most promising option., Perhaps spatial averaging of these equations to obtain a one (or two zone) models like as done by \citet{Pac83} may be the most promising option.
+ Such approximations may be obtained with insights gained from the results of future sophisticated numerical magneto-hydrodynamic simulations of realistic accretion disks., Such approximations may be obtained with insights gained from the results of future sophisticated numerical magneto-hydrodynamic simulations of realistic accretion disks.
+ Non-linear time series analysis can also be used to compare results of numerical simulations with real data., Non-linear time series analysis can also be used to compare results of numerical simulations with real data.
+" If the simulation correctly encompasses the basic non-linear behavior of the system, the resultant simulated light curve should have similar D?(M) values and SVD reconstruction."," If the simulation correctly encompasses the basic non-linear behavior of the system, the resultant simulated light curve should have similar $D_2 (M)$ values and SVD reconstruction."
+" In analogy with the Lorenz system, the different temporal behavior (corresponding to the different classes) can be attributed to variation of a single control parameter, which for the Lorenz system is R (Eqn 1))."," In analogy with the Lorenz system, the different temporal behavior (corresponding to the different classes) can be attributed to variation of a single control parameter, which for the Lorenz system is $R$ (Eqn \ref{loreqn}) )."
+" For GRS 1915-105, this controlling parameter could be the time averaged accretion rate, which is determined by the outer boundary condition."," For GRS 1915+105, this controlling parameter could be the time averaged accretion rate, which is determined by the outer boundary condition."
+ One of the key attributes of a deterministic non-linear system is that the time variability need not be due to a time varying parameter., One of the key attributes of a deterministic non-linear system is that the time variability need not be due to a time varying parameter.
+" For example, the equations defining the Lorenz system (Eqn 1)) do not have an explicit time term."," For example, the equations defining the Lorenz system (Eqn \ref{loreqn}) ) do not have an explicit time dependent term."
+" For GRS 1915-4105, it is possible that thedependent different temporal classes arise from variations in the accretion rate, while the variability for a given class is due to an inherent non-"," For GRS 1915+105, it is possible that the different temporal classes arise from variations in the accretion rate, while the variability for a given class is due to an inherent non-linearity."
+" Thus it is not necessary to postulate other varying parameters, apart from the accretion rate (like magnetic field etc) to explain the source's complex variability."," Thus it is not necessary to postulate other varying parameters, apart from the accretion rate (like magnetic field etc) to explain the source's complex variability."
+" If this is true, then the average bolometric luminosity which should be proportional to the average accretion rate, should be different for each class."," If this is true, then the average bolometric luminosity which should be proportional to the average accretion rate, should be different for each class."
+" However, the bolometric luminosity can be estimated only by model dependent spectral fitting with the caveat that the time-averaged may not be adequately represented by steady ones."," However, the bolometric luminosity can be estimated only by model dependent spectral fitting with the caveat that the time-averaged spectra may not be adequately represented by steady ones."
+" spectraNevertheless, the results presented here provide the possibility of identifying a single control parameter which determines the temporal class exhibited by the system."," Nevertheless, the results presented here provide the possibility of identifying a single control parameter which determines the temporal class exhibited by the system."
+" The temporal behavior of GRS 19154105 on time-scales of hours and months is unique, but may be related to the state changes that occur for other black hole systems (like Cygnus X-1) on time-scales of months."," The temporal behavior of GRS 1915+105 on time-scales of hours and months is unique, but may be related to the state changes that occur for other black hole systems (like Cygnus X-1) on time-scales of months."
+ This difference between GRS 19154105 and the other black holes has led to the speculation that GRS 19154105 has rapidly spinning black hole., This difference between GRS 1915+105 and the other black holes has led to the speculation that GRS 1915+105 has a rapidly spinning black hole.
+" the non-lineara equations for GRS 19154105 will Identifyingnot only governingprovide insight into the temporal variability of other black hole systems, but will perhaps give an explanation for its uniqueness."," Identifying the governing non-linear equations for GRS 1915+105 will not only provide insight into the temporal variability of other black hole systems, but will perhaps give an explanation for its uniqueness."
+for determining the amplitudes of toroidal and poloidal fields. aid hence the polar fields.,"for determining the amplitudes of toroidal and poloidal fields, and hence the polar fields."
+ The primary role of meridional circulation in f[lux-transport dvnamo models is the advective transport of magnetic fields. and hence the structure and strength of the flow are crucial in the dvnamo models for determining the timings. namely the duration of a evcele. ils rise and fall pattern ancl the timing of the reversals of the Suns polar fields (Dikpati&2004:Dikpatietal 2010).," The primary role of meridional circulation in flux-transport dynamo models is the advective transport of magnetic fields, and hence the structure and strength of the flow are crucial in the dynamo models for determining the timings, namely the duration of a cycle, its rise and fall pattern and the timing of the reversals of the Sun's polar fields \citep{dc99,dikpati2004,dgdu2010}."
+.. A transport process like meridional circulation in a fIux-transport dvnamo redistributes the dvnamo-generated magnetic flux. a relatively minor effect in creating an increase or decrease of magnetic flux.," A transport process like meridional circulation in a flux-transport dynamo redistributes the dynamo-generated magnetic flux, a relatively minor effect in creating an increase or decrease of magnetic flux."
+ This is illustrated in Table 1 of Dikpati&Charbonneau(1999) in which it is seen (hat. if the meridional circulation speed is doubled. the peak polar field changes by only 354. whereas for the same meridional circulation. increasing (he surface poloidal source by a factor 2.5 doubles the peak polar field. and decreasing the surface source by 75% decreases the polar field peak by 89%.," This is illustrated in Table 1 of \citet{dc99} in which it is seen that, if the meridional circulation speed is doubled, the peak polar field changes by only $3\%$, whereas for the same meridional circulation, increasing the surface poloidal source by a factor 2.5 doubles the peak polar field, and decreasing the surface source by $75\%$ decreases the polar field peak by $89\%$."
+ The question we address is how flux-transport dvnanmo models and surface (ransport models respond to (he changes in (he surface poleward flow. speed., The question we address is how flux-transport dynamo models and surface transport models respond to the changes in the surface poleward flow speed.
+ Detailed calculations show that a surface transport model produces a weaker polar field (Wane.Lean&SheelevSchrijver&Liu 2008).. while a flux transport dvnanmo gives a stronger polar field when (he poleward surface flow-speed is increased (Dikpati&Charbonnean1999:Dikpati.deToma&Gilman 2008).," Detailed calculations show that a surface transport model produces a weaker polar field \citep{wls2005, sl2008}, while a flux transport dynamo gives a stronger polar field when the poleward surface flow-speed is increased \citep{dc99,ddg2008}."
+. Are these opposite results in conflict?, Are these opposite results in conflict?
+ We address. first. with qualitative reasoning and schematics and then in the next section with numerical solution from a fhix-transport dvnamo model. whether (his is a true conflict. ancl if so. what is the physical origin of (his eondlict.," We address, first with qualitative reasoning and schematics and then in the next section with numerical solution from a flux-transport dynamo model, whether this is a true conflict, and if so, what is the physical origin of this conflict."
+ Figure | ilustrates three scenarios: (a) and (b) for surface flix (ransport models aud (c) for a flux transport dvnamo., Figure 1 illustrates three scenarios: (a) and (b) for surface flux transport models and (c) for a flux transport dynamo.
+ The primary difference between figures 1(a) and (b) is the latitude where the surface polewarel flow is maxinum., The primary difference between figures 1(a) and (b) is the latitude where the surface poleward flow is maximum.
+" Figure 1(a) illustrates what happens to surface [Iux in the decliningphase of a evele when the meridional [Iow peaks al 6"" as in ihe model by Wang.Lean&Sheeley(2005): Figure 1(b) illustrates what happens for a meridional [low (hat peaks al 37.5° (Schrijver.DeRosa&Title2002).", Figure 1(a) illustrates what happens to surface flux in the decliningphase of a cycle when the meridional flow peaks at $6^{\circ}$ as in the model by \citet{wls2005}; Figure 1(b) illustrates what happens for a meridional flow that peaks at $37.5^{\circ}$ \citep{sdt2002}.
+. Figure 1(c) shows what happens to polar fields when a merilional circulation profile in the pole-to-equator meridional plane peaks at mid-latitudes ~ 40°. as has been used in fIux-transport dvnamo models by Dikpati&Charbonneau(1999) and Dikpati.deToma&Gilman (2005).," Figure 1(c) shows what happens to polar fields when a meridional circulation profile in the pole-to-equator meridional plane peaks at mid-latitudes $\sim 40^{\circ}$ , as has been used in flux-transport dynamo models by \citet{dc99} and \citet{ddg2008}. ."
+han as a thermally ciriven [ow (Le. with wyο) and i.~pgO where den is the Ekman-Lartmann boundary-laver hickness).,than as a thermally driven flow (i.e. with $\ut \sim \rc\tilde{\Omega}$ and $\ur \sim \delta_{\rm EH} \tilde{\Omega}$ where $\delta_{\rm EH}$ is the Ekman-Hartmann boundary-layer thickness).
+ “Phis discrepancy. should. influence quantitative owedietions of the model only whilst. leaving qualitative results unalfected., This discrepancy should influence quantitative predictions of the model only whilst leaving qualitative results unaffected.
+ One of the main aims of this study is to be able to oedicet the angular velocity of the interior., One of the main aims of this study is to be able to predict the angular velocity of the interior.
+" In order to do so. the angular velocity of the inner core Qj, is treated as an eigenvalue of the problem. ancl an additional boundary condition is imposed on the svstem accordingly. namely that no net torque is applied. to the region interior tor =ra (this is à necessary condition to guarantee à steady state)."," In order to do so, the angular velocity of the inner core $\Omega_{\rm in}$ is treated as an eigenvalue of the problem, and an additional boundary condition is imposed on the system accordingly, namely that no net torque is applied to the region interior to $r=\rin$ (this is a necessary condition to guarantee a steady state)."
+" This condition. which determines uniquely the value of Qin. is equivalent to requiring that the integral of the momentum flux through the boundary. vanishes: One can notice through this expression the main justification for choosing conducting boundary. conclitions: when the region below rj, is insulating. the racial component of the current (and thereby the toroidal field) must. vanish on the boundarv: in that case the angular-momentunm Hux through the boundary would be purely viscous."," This condition, which determines uniquely the value of $\Omega_{\rm in}$, is equivalent to requiring that the integral of the angular-momentum flux through the boundary vanishes: One can notice through this expression the main justification for choosing conducting boundary conditions: when the region below $\rin$ is insulating, the radial component of the current (and thereby the toroidal field) must vanish on the boundary; in that case the angular-momentum flux through the boundary would be purely viscous."
+ Since viscous stresses are normally thought to be negligible in the sun. such a model would be a poor representation of the dynamical structure of the radiative zone.," Since viscous stresses are normally thought to be negligible in the sun, such a model would be a poor representation of the dynamical structure of the radiative zone."
+ The system of equations (5)) and the boundary conditions presented in Section 2.3. constitute a wellkposecl partial differential equations svstem with one eigenvalue., The system of equations \ref{eq:eqsmhd}) ) and the boundary conditions presented in Section \ref{sec:bc} constitute a well-posed partial differential equations system with one eigenvalue.
+ The numerical method chosen for the solution of this svsteni is the following: A more detailed description of the numerical procedure can be found in the work by Garaucl (2001b)., The numerical method chosen for the solution of this system is the following: A more detailed description of the numerical procedure can be found in the work by Garaud (2001b).
+ Alanual mesh-stretehing is preferred. over. automatic mesh-point allocation. as the latter was found to have low performance for. bounclary-laver thicknesses order of 10.μαMn) or less.," Manual mesh-stretching is preferred over automatic mesh-point allocation, as the latter was found to have low performance for boundary-layer thicknesses order of $10^{-6}
+(\rout-\rin)$ or less."
+ This low performance is probably due to the fact that as many Fourier modes απο used. the allocation algorithms that were tried. cannot choose adequately according the most relevant function against which the mesh shoulcl be stretched.," This low performance is probably due to the fact that as many Fourier modes are used, the allocation algorithms that were tried cannot choose adequately according the most relevant function against which the mesh should be stretched."
+ When stretching the mesh manually (according to the width anc positions of the boundary layers determined in Section 5)). a minimum. of a hundred. points is allocated to cach of the boundary lavers.," When stretching the mesh manually (according to the width and positions of the boundary layers determined in Section \ref{sec:blanal}) ), a minimum of a hundred points is allocated to each of the boundary layers."
+ “Pvpically. another 200-400 points are allocated to the interval outside the boundary lavers. depending on the simulations.," Typically, another 200-400 points are allocated to the interval outside the boundary layers, depending on the simulations."
+ As a first step towards the resolution of the model presented above. a simplified system is solved in which the magnetic Ποια is ignored.," As a first step towards the resolution of the model presented above, a simplified system is solved in which the magnetic field is ignored."
+ Phe hvdrodynamies of a Buid between two concentric impermeable rotating spheres is a relatively well-studied problem. both analytically (since Proudman. 1956) and numerically (see for instance Dormy. Cardin Jault. 1998): these previous studies can be compared with the results of the non-magnetic numerical solutions obtained with the numerical procedure presented in Section 2.4. to test its accuracy and performance.," The hydrodynamics of a fluid between two concentric impermeable rotating spheres is a relatively well-studied problem, both analytically (since Proudman, 1956) and numerically (see for instance Dormy, Cardin Jault, 1998); these previous studies can be compared with the results of the non-magnetic numerical solutions obtained with the numerical procedure presented in Section \ref{sec:num} to test its accuracy and performance."
+" As a example. the solution to the non-magnetic subset of equations (5)) with an Ekman number ££=8.10"" is presented in Fig. 2.."," As a example, the solution to the non-magnetic subset of equations \ref{eq:eqsmhd}) ) with an Ekman number $\Enu = 8 \times 10^{-6}$ is presented in Fig. \ref{fig:nomagn}."
+ When no magnetic field is present. for low enough Ekman number. Duid. motion is dominated. by Coriolis forces everywhere except in two boundary. lavers near the spherical boundaries. and in a shear laver at the tangent evlinder.," When no magnetic field is present, for low enough Ekman number, fluid motion is dominated by Coriolis forces everywhere except in two boundary layers near the spherical boundaries, and in a shear layer at the tangent cylinder."
+ In the bulk. of the Ελιά. angular velocity is more-or-less constant on cvlinders: indeed. when Viscosity is negligible. the fluid dynamics equations for the incompressible Εις reduce to which implies that 2£ must be parallel to the rotation axis. ancl which implies that the angular velocity must. be independent of z where z is the evlindrical coordinate that runs parallel to the rotation axis (Proudman. 1916. Taylor. 1921).," In the bulk of the fluid, angular velocity is more-or-less constant on cylinders: indeed, when viscosity is negligible, the fluid dynamics equations for the incompressible fluid reduce to which implies that $\bu$ must be parallel to the rotation axis, and which implies that the angular velocity must be independent of $z$ where $z$ is the cylindrical coordinate that runs parallel to the rotation axis (Proudman, 1916, Taylor, 1921)."
+ Viscous elfects are necessary in the boundary. lavers to ensure the smooth transition between the rotation profile in the bulk of the [uid and that imposed at the boundaries., Viscous effects are necessary in the boundary layers to ensure the smooth transition between the rotation profile in the bulk of the fluid and that imposed at the boundaries.
+" ""his structure was first studied by Proudman (1956) and Stewartson (1966). in the case where ££, is asvmptotically small."," This structure was first studied by Proudman (1956) and Stewartson (1966), in the case where $E_{\nu}$ is asymptotically small."
+ It is possible to show (see the Appendix) that in this limit the interior angular velocity is uniquely determined by the size of the gap between the two spheres: its value can be predicted analytically. and is given by the following equation:," It is possible to show (see the Appendix) that in this limit the interior angular velocity is uniquely determined by the size of the gap between the two spheres; its value can be predicted analytically, and is given by the following equation:"
+density of 500 XBPs for the entire solar surface at any given time.,density of 800 XBPs for the entire solar surface at any given time.
+ It is known that the observed XBP number is anti-correlated with the solar cycle. but this is an observational bias and the umber density of XBPs ts nearly independent of the 11-yr solar activity eycle (Nakakubo Hara 1999; Sattarov et al.," It is known that the observed XBP number is anti-correlated with the solar cycle, but this is an observational bias and the number density of XBPs is nearly independent of the 11-yr solar activity cycle (Nakakubo Hara 1999; Sattarov et al."
+ 2002: Hara Nakakubo 2003)., 2002; Hara Nakakubo 2003).
+ Golub et al. (, Golub et al. (
+"1974) found that the ""Siameters of the XBPs are around 10-20 are sec and their life times range from 2 hours to 2 days (Zhang et al.",1974) found that the diameters of the XBPs are around 10-20 arc sec and their life times range from 2 hours to 2 days (Zhang et al.
+ 2001)., 2001).
+ The physical studies have indicated the temperatures to be fairly low. T=210° K. and the electron densities = 5 xIO0*cur? (Golub Pasachoff 1997).," The physical studies have indicated the temperatures to be fairly low, $T= 2 \times 10^6$ K, and the electron densities = 5 $\times 10^{9} cm^{-3}$ (Golub Pasachoff 1997)."
+ Recently. Kartyappa and Varghese (2008) have analysed the time sequence of XBPs in soft X-ray images obtained from Hinode/XRT and discussed the nature of intensity oscillations associated with the different classes of XBPs showing the different emission levels.," Recently, Kariyappa and Varghese (2008) have analysed the time sequence of XBPs in soft X-ray images obtained from Hinode/XRT and discussed the nature of intensity oscillations associated with the different classes of XBPs showing the different emission levels."
+ It has been shown by various groups that the solar corona rotates differentially on the basis of their analysis of coronal bright points in SOHO/EIT images (Brajsa et al., It has been shown by various groups that the solar corona rotates differentially on the basis of their analysis of coronal bright points in SOHO/EIT images (Brajsa et al.
+ 2001. 2002. 2004: Karachik et al.," 2001, 2002, 2004; Karachik et al."
+ 2006). and by using radio observations at different frequencies (Vats et al.," 2006), and by using radio observations at different frequencies (Vats et al."
+ 2001)., 2001).
+ However. further investigations are required to confirm the differential rotation of the corona.," However, further investigations are required to confirm the differential rotation of the corona."
+ Particularly. it is not clear that the sidereal angular rotation velocity of the corona depends on the sizes and lifetimes of the tracers (XBPs) and how it varys with the phases of the solar cycle.," Particularly, it is not clear that the sidereal angular rotation velocity of the corona depends on the sizes and lifetimes of the tracers (XBPs) and how it varys with the phases of the solar cycle."
+ In this paper. we have analysed two sets of observations of full-disk soft X-ray images obtained from Hinode/XRT and Yohkoh/SXT experiments to determine the coronal rotation using the XBPs as tracers and explore the evidence for differential rotation and its relation to the sizes of the XBPs and phases of the solar magnetic cycle.," 	 In this paper, we have analysed two sets of observations of full-disk soft X-ray images obtained from Hinode/XRT and Yohkoh/SXT experiments to determine the coronal rotation using the XBPs as tracers and explore the evidence for differential rotation and its relation to the sizes of the XBPs and phases of the solar magnetic cycle."
+ We have used daily full-disk soft X-ray images obtained during the period of January. March and April. 2007 from X-Ray Telescope (XRT) on-board the Hinode mission.," We have used daily full-disk soft X-ray images obtained during the period of January, March and April, 2007 from X-Ray Telescope (XRT) on-board the Hinode mission."
+" The images have been observed through a single X-ray ppoly filter and the image size is 20487x2048"" with a spatial resolution of 1.0327/pixel."," The images have been observed through a single X-ray poly filter and the image size is 2048""x2048"" with a spatial resolution of 1.032""/pixel."
+ All the full-disk X-ray images were visually inspected and the XBPs have been identified in sequences of images., All the full-disk X-ray images were visually inspected and the XBPs have been identified in sequences of images.
+ The main criterion for the identification of a XBP was checking its persistence 1n the consecutive Images at approximately the same latitude and shifted in the Central Meridian Distance (CMD) according to the elapsed time., The main criterion for the identification of a XBP was checking its persistence in the consecutive images at approximately the same latitude and shifted in the Central Meridian Distance (CMD) according to the elapsed time.
+ CMD values of the identified features were then measured in selected consecutive images and were fitted as a function of time., CMD values of the identified features were then measured in selected consecutive images and were fitted as a function of time.
+ The correlation coetficient of the function CMD was generally very high. implying that the tracers were correctly identified.," The correlation coefficient of the function CMD was generally very high, implying that the tracers were correctly identified."
+ We have identified and selected a well isolated and distinct 63 XBPs which are distributed over the different heliographic latitudes., We have identified and selected a well isolated and distinct 63 XBPs which are distributed over the different heliographic latitudes.
+ Using SSW in IDL. we have generated the full-disk maps of the images and overlayed them with longitude and latitude grid maps.," Using SSW in IDL, we have generated the full-disk maps of the images and overlayed them with longitude and latitude grid maps."
+ Using the tracer method (the tracer method is based on a visual identification of a particular XBP that can be used as a tracer). coronal X-ray bright points," Using the tracer method (the tracer method is based on a visual identification of a particular XBP that can be used as a tracer), coronal X-ray bright points"
+mixing angle aud Dirac type CP-violating phase. 0. are given for these textures.,"mixing angle and Dirac type CP-violating phase, $\delta$, are given for these textures."
+ These {wo parameters are expected (o be measured in the forthcoming neutrino oscillation experimentis., These two parameters are expected to be measured in the forthcoming neutrino oscillation experiments.
+ We. also. obtained the lower bound on effective Majorana mass for all the allowed hybrid textures.," We, also, obtained the lower bound on effective Majorana mass for all the allowed hybrid textures."
+ The possible measurement of effective Majorana mass in neutrinoless double ;2 decay experiments will provide an additional constraint on the remaining three neutrino parameters i.e. (he neutiino mass scale and (wo Majorana ivpe CP-violating phases., The possible measurement of effective Majorana mass in neutrinoless double $\beta$ decay experiments will provide an additional constraint on the remaining three neutrino parameters i.e. the neutrino mass scale and two Majorana type CP-violating phases.
+" The observation of correlations between. various neutrino parameters like 044. 054. 0. M, ele will confirm/reject hybrid textures with a texture zero and an additional equality among the matrix elements."," The observation of correlations between various neutrino parameters like $\theta_{13}$ , $\theta_{23}$, $\delta$, $M_{ee}$ etc will confirm/reject hybrid textures with a texture zero and an additional equality among the matrix elements."
+in the range 10* to several LO? AL.pe. and NV = several 107-10? (see. e.g.. Pfalzner 2009).,"in the range $10^{3}$ to several $^{5}$ $M_{\odot} \ pc^{-3}$ and $N$ = several $^4$ $^5$ (see, e.g., Pfalzner 2009)."
+ Similar trends have been found for star clusters in other galaxies (Plalzner Eckart 2009)., Similar trends have been found for star clusters in other galaxies (Pfalzner Eckart 2009).
+ The above results can only be properly understood within the context of the Galactic tidal field: i.e. cluster Udal radi and separations plav a major role in (he outcome οἱ primordial binary cluster evolution., The above results can only be properly understood within the context of the Galactic tidal field; i.e. cluster tidal radii and separations play a major role in the outcome of primordial binary cluster evolution.
+ The strength of the cIuster-cluster interaction is maximum when the intercIuster separation becomes smaller Chan the tidal radii., The strength of the cluster-cluster interaction is maximum when the intercluster separation becomes smaller than the tidal radii.
+ This interpretation was proposed in de la Fuente Marcos de la Fuente Marcos (20000) in the form of a classification scheme (8)) based only in the values of separation GS) and tidal radii (Hj. / = 1. 2): 'The scheme was implemented by considering the available observational evidence.," This interpretation was proposed in de la Fuente Marcos de la Fuente Marcos (2009c) in the form of a classification scheme \ref{bincri}) ) based only in the values of separation $S$ ) and tidal radii $R_{Ti}$, $i$ = 1, 2): The scheme was implemented by considering the available observational evidence."
+ For open Cluster pairs of similar mass ratio in the detached and weakly interacting categories. ionization into the separated (wins state is the observed. evolutionary. path.," For open cluster pairs of similar mass ratio in the detached and weakly interacting categories, ionization into the separated twins state is the observed evolutionary path."
+ Semi-detached aud. pairs merge in a Gmescale that depends strongly on the orbital eccentricity., Semi-detached and in-contact pairs merge in a timescale that depends strongly on the orbital eccentricity.
+ Very eccentric pairs merge in a timescale of nearly 10 Myr ancl. therefore. the actual observation of (he merging process may be very dillicult as it may happen even before the embedded phase ends.," Very eccentric pairs merge in a timescale of nearly 10 Myr and, therefore, the actual observation of the merging process may be very difficult as it may happen even before the embedded phase ends."
+ If (he mass ratio is appreciably different. extreme tidal distortion lollowed by actual destruction of the smallest. cluster is always observed.," If the mass ratio is appreciably different, extreme tidal distortion followed by actual destruction of the smallest cluster is always observed."
+ The above interpretation has strong implications on what should be observed within the Milkv. Way al different Galactocentric distances ancl in other galaxies., The above interpretation has strong implications on what should be observed within the Milky Way at different Galactocentric distances and in other galaxies.
+ Weaker ealactic (dal fields increase the probability of observing semi-detached or in-contact cluster pairs and. eventually. mergers.," Weaker galactic tidal fields increase the probability of observing semi-detached or in-contact cluster pairs and, eventually, mergers."
+ The tidal force eradient determines the cluster tidal radius (e.g. Elmegreen IIunter 2010)., The tidal force gradient determines the cluster tidal radius (e.g. Elmegreen Hunter 2010).
+ Even if the fraction of primordial binary clusters may well be similar across different galaxies. the preferential evolutionary. patli could be rather different: ionization being dominant in massive ealaxies and the central regions of galaxies in general.," Even if the fraction of primordial binary clusters may well be similar across different galaxies, the preferential evolutionary path could be rather different: ionization being dominant in massive galaxies and the central regions of galaxies in general."
+ Figures l.. 4.. 6. and 7.. also show probable primordial pairs (age difference «30 Myr) with pair age < 210 Myr.," Figures \ref{evolequal}, \ref{evolhalfsd}, \ref{evolhalfdd} and \ref{evolfourth}, also show probable primordial pairs (age difference $<$ 30 Myr) with pair age $<$ 210 Myr."
+ Following de la Fuente Marcos de la Fuente Marcos. (2009b).," Following de la Fuente Marcos de la Fuente Marcos (2009b),"
+"After a WD in the SDSS catalog is identified by aas a member of a binary system, follow-up observations are necessary to determine its period and amplitude, and obtain high signal-to-noise spectra that can be used to measure its effective temperature Typ, and gravity logg, calculate its mass, and constrain the nature of its companion.","After a WD in the SDSS catalog is identified by as a member of a binary system, follow-up observations are necessary to determine its period and amplitude, and obtain high signal-to-noise spectra that can be used to measure its effective temperature $T_{eff}$ and gravity $\log{g}$, calculate its mass, and constrain the nature of its companion."
+" In the case of12574-5428,, the follow-up observations were conducted with the Dual Imaging Spectrograph (DIS) at the 3.5m ARC telescope at APO.", In the case of the follow-up observations were conducted with the Dual Imaging Spectrograph (DIS) at the 3.5m ARC telescope at APO.
+" A total of 23 spectra were taken on the nights of February 5, 6, and 14, 2009."," A total of 23 spectra were taken on the nights of February 5, 6, and 14, 2009."
+" We used the B1200 grating with a 1.5"" slit to achieve a 0.62 ddispersion and a FWHM 1.8 rresolution."," We used the B1200 grating with a 1.5"" slit to achieve a 0.62 dispersion and a FWHM 1.8 resolution."
+" The integration time was 10 minutes for all spectra, bracketed by exposures of He, Ne and Ar arc lamps to ensure a proper wavelength calibration."," The integration time was 10 minutes for all spectra, bracketed by exposures of He, Ne and Ar arc lamps to ensure a proper wavelength calibration."
+ Flux was calibrated each night by taking several exposures of the spectrophotometric standard Feige 67 (Oke1990).., Flux was calibrated each night by taking several exposures of the spectrophotometric standard Feige 67 \citep{oke90:spectrophotometric_standards}.
+ The data were reduced using standard long-slit IRAF routines (Tody1993).., The data were reduced using standard long-slit IRAF routines \citep{tody93:IRAF}.
+" A prominent Hg emission line at 4358Á,, courtesy of the inhabitants of White Sands, NM, confirmed that the wavelengths are calibrated to better than the instrumental resolution: the standard deviation of the line over 8 spectra (one night) is 0.1A."," A prominent Hg emission line at 4358, courtesy of the inhabitants of White Sands, NM, confirmed that the wavelengths are calibrated to better than the instrumental resolution: the standard deviation of the line over 8 spectra (one night) is 0.1."
+". The final spectra span between 3770 and 5030Á,, and clearly show all the lines in the Balmer series from Hg to Hio."," The final spectra span between 3770 and 5030, and clearly show all the lines in the Balmer series from $_{\beta}$ to $_{10}$."
+ We measured the RV shift of each APO spectrum by fitting the centroid of the line after applying the standard Solar System barycentricHg corrections., We measured the RV shift of each APO spectrum by fitting the centroid of the $_{\beta}$ line after applying the standard Solar System barycentric corrections.
+" We used a Lorentzian profile with a Gaussian core (Thompsonetal.2004),, adjusting the fit parameters with MPFIT, a version of the Levenberg-Marquardt algorithm adapted for IDL (Markwardt2009).."," We used a Lorentzian profile with a Gaussian core \citep{thompson04:PY_Vul}, adjusting the fit parameters with MPFIT, a version of the Levenberg-Marquardt algorithm adapted for IDL \citep{markwardt09:MPFIT}."
+" The resulting RV curve is well fit by a circular orbit of the form K4sin((2x(t—To)/P)) where P is the period, K4 the semiamplitude, 7o a temporal--yA, offset marking the zero-point of the curve, and γα a constant RV offset (see Figure 2))."," The resulting RV curve is well fit by a circular orbit of the form $K_{A}\sin((2 \pi (t-T_{0})/P))+\gamma_{A}$, where $P$ is the period, $K_{A}$ the semiamplitude, $T_{0}$ a temporal offset marking the zero-point of the curve, and $\gamma_{A}$ a constant RV offset (see Figure \ref{fig-2}) )."
+" We found that the nine day baseline of theAPO spectra strongly constrains the period of the binary to P=4.5550+0.0007 hr, (Figure 3)), with a semiamplitude of K4=322.7--6.3kms! (see Table 1 for a complete list of the orbital parameters)."," We found that the nine day baseline of theAPO spectra strongly constrains the period of the binary to $P=4.5550 \pm 0.0007$ hr, (Figure \ref{fig-3}) ), with a semiamplitude of $K_{A}=322.7 \pm 6.3\,\mathrm{km\,s^{-1}}$ (see Table \ref{tab-1} for a complete list of the orbital parameters)."
+" The results from E06 give Te=8750+25 K and logg=9.0+0.005 for 1257+5428,,but"," The results from E06 give $T_{eff}=8750\pm25$ K and $\log{g}=9.0\pm0.005$ for ,but"
+the other cells in this cube).,the other cells in this cube).
+ Third. ρω. a measure of the density surrounding this cube. defined as All eight cells used to find pc are marked as active if The reasons for such a choice of form can be seen by examining the behavior of eq. [5]]," Third, $\rho_{\rm ext}$, a measure of the density surrounding this cube, defined as All eight cells used to find $\rho_c$ are marked as active if The reasons for such a choice of form can be seen by examining the behavior of eq. \ref{eqn:rhocrit}] ]"
+ in various limits., in various limits.
+ At early times in a cosmological simulation p2 and o:«p. 5ο pipe7$Apert~Αρ.," At early times in a cosmological simulation $\rho \approx \bar{\rho}$ and $\sigma\ll \bar{\rho}$, so $\rho_{\rm thr} \approx \onehalf A\rho_{\rm ext} \approx \onehalf A\bar{\rho}$ ."
+ Thus with a choice of A=2 those regions which are only slightly overdense will be selected., Thus with a choice of $A=2$ those regions which are only slightly overdense will be selected.
+ At late times 3p., At late times $\sigma\gg \bar{\rho}$.
+ In void regious peaKo. which means pyyy72Ppeg: thus an isolated cell will be chosen if it is a fixed multiple above its surroundings.," In void regions $\rho_{\rm ext}\ll \sigma$, which means $\rho_{\rm thr} \approx B\rho_{\rm ext}$; thus an isolated cell will be chosen if it is a fixed multiple above its surroundings."
+" Lon the other hand poy2o. then piu,&Bo. meaning that at each step there is a given density above which all cells will be ehoseu: this limit will increase with time as o increases."," If on the other hand $\rho_{\rm ext}\gg \sigma$, then $\rho_{\rm thr} \approx B\sigma$, meaning that at each step there is a given density above which all cells will be chosen; this limit will increase with time as $\sigma$ increases."
+ Ixeep in mind that the distribution of cell deusities will be approximately lognormal (Ixayo.Taruya&Suto(2001) aud references therein) when interpreting the value of o., Keep in mind that the distribution of cell densities will be approximately lognormal \citet{KayoTaSu01} and references therein) when interpreting the value of $\sigma$.
+ Alter cell selection. the domain decomposition proceeds in the manner described inBOX.," After cell selection, the domain decomposition proceeds in the manner described in."
+ Cells are linked by friends-of-Driends. vielding regions of space separated by at least one cell length.," Cells are linked by friends–of–friends, yielding regions of space separated by at least one cell length."
+ Any particle which contributes some mass to ani active cell when finding the PM density is assigued to tlie correspoudiug tree., Any particle which contributes some mass to an active cell when finding the PM density is assigned to the corresponding tree.
+ Because tlie volume of all tree regions[we is less than oue percent ofthe total volume. the amount oL memory allocated to hold the tidal potential data is generally not siguificaut compared to tliat already needed for the PAL mesh.," Because the volume of all tree regions is less than one percent of the total volume, the amount of memory allocated to hold the tidal potential data is generally not significant compared to that already needed for the PM mesh."
+ However. oue complication of this scheme is that at early times. wheu o is suiall. the resulting trees can be very long filaments.," However, one complication of this scheme is that at early times, when $\sigma$ is small, the resulting trees can be very long filaments."
+ While not coutaiuing a large amount of matter. beiug not very overdeuse and nearly one-dimiensioual. the size of a cubical volume euclosing such a tree will sometimes be a significant fraction of the entire simulation volume.," While not containing a large amount of matter, being not very overdense and nearly one–dimensional, the size of a cubical volume enclosing such a tree will sometimes be a significant fraction of the entire simulation volume."
+ This can cause dilliculties because of the amount ofcomputer memory required to compute and save the tidal potential in such subvolumes., This can cause difficulties because of the amount ofcomputer memory required to compute and save the tidal potential in such subvolumes.
+ The value of B in eq. [5]], The value of $B$ in eq. \ref{eqn:rhocrit}] ]
+ is thus allowed to increase at early times when the spatial extent of trees teuds to become larger. aud decrease at later times when trees are more compact.," is thus allowed to increase at early times when the spatial extent of trees tends to become larger, and decrease at later times when trees are more compact."
+ We have settled on the following method. which has worked well iu a variety of simulatious.," We have settled on the following method, which has worked well in a variety of simulations."
+ Space is allocated for a maximum subvolume leneth oue quarter of the PM mesh size., Space is allocated for a maximum subvolume length one quarter of the PM mesh size.
+ The size of the largest sub-box used is cliecked at the enc of each PM step., The size of the largest sub-box used is checked at the end of each PM step.
+ At early times (when o« 1.6) B is increased by if tliis box has a leneth more than a third of the allocated value., At early times (when $\sigma<1.6$ ) $B$ is increased by if this box has a length more than a third of the allocated value.
+ Ouce o exceeds 1.6 iu a typical simulation.E. the spatial extent of trees has stopped growing. so iu this case B is decreased by if the largest sub-box has a leneth less than half the maximum allocated.," Once $\sigma$ exceeds 1.6 in a typical simulation, the spatial extent of trees has stopped growing, so in this case $B$ is decreased by if the largest sub-box has a length less than half the maximum allocated."
+ A minimum value is set for B to prevent too mauy tree particles beiug chosen: a choice of B=10 will place roughlyhalf of the particles in trees at 2= 0., A minimum value is set for $B$ to prevent too many tree particles being chosen; a choice of $B=10$ will place roughlyhalf of the particles in trees at $z=0$ .
+ An example of this in practice is cliscussed in refsec:overv.., An example of this in practice is discussed in \\ref{sec:overv}. .
+"When undergoing an outburst, BHC sources usually follow a relatively well known “q-track” pattern in the HID 222).","When undergoing an outburst, BHC sources usually follow a relatively well known “q-track” pattern in the HID ."
+" The outburst starts in the so-called low-hard spectral state (LHS), which is characterized by a power-law shaped X-ray spectrum with [~1.4-1.5 and a cut-off at the higher energies (Εωι-100 ΚΚον)."," The outburst starts in the so-called low-hard spectral state (LHS), which is characterized by a power-law shaped X-ray spectrum with $\Gamma$$\simeq$ 1.4-1.5 and a cut-off at the higher energies $E_{\rm cut}$$\sim$ keV)."
+ A soft thermal component with kT<0.5 kkeV is often observed in the LHS and ascribed to the emission from an accretion disk truncated at long distances from the central BH?)., A soft thermal component with $kT<0.5$ keV is often observed in the LHS and ascribed to the emission from an accretion disk truncated at long distances from the central BH.
+. The radio emission in this state is caused by the synchrotron radiation of a steady jet., The radio emission in this state is caused by the synchrotron radiation of a steady jet.
+" During the early rise of the outburst, the X-ray and radio luminosities both increase, but the X-ray color of the spectrum remains hard(??)."," During the early rise of the outburst, the X-ray and radio luminosities both increase, but the X–ray color of the spectrum remains hard."
+". As the outburst progresses, the source reaches the high-soft state (HSS) characterized by a prominent thermal emission from the accretion disk and a marginal power law tail."," As the outburst progresses, the source reaches the high-soft state (HSS) characterized by a prominent thermal emission from the accretion disk and a marginal power law tail."
+" In addition to these two main states, identify the hard-intermediate (HIMS) and soft-intermediate (SIMS) states with spectral parameters of the disk and power-law component in between the LHS and HSS."," In addition to these two main states, identify the hard-intermediate (HIMS) and soft-intermediate (SIMS) states with spectral parameters of the disk and power-law component in between the LHS and HSS."
+ We will exploit this classification in the remainder of this discussion., We will exploit this classification in the remainder of this discussion.
+" An alternative division of the source intermediate states was used, e.g., by?,, who define a steep power-law state (I'> 2.4) generally observed at high flux and then also known as very high state, plus an intermediate state that covers the unclassified observations."," An alternative division of the source intermediate states was used, e.g., by, who define a steep power-law state $\Gamma>2.4$ ) generally observed at high flux and then also known as very high state, plus an intermediate state that covers the unclassified observations."
+ We stress that both SIMS and HSS are characterized by the absence of radio emission., We stress that both SIMS and HSS are characterized by the absence of radio emission.
+ This is interpreted as the suppression of the jet and marks the most striking difference between the soft and hard states., This is interpreted as the suppression of the jet and marks the most striking difference between the soft and hard states.
+ The short timescale variability also varies during the outbursts., The short timescale variability also varies during the outbursts.
+ The level of the RMS noise decreases with increasing flux in the LHS and decreases even more when leaving this state., The level of the RMS noise decreases with increasing flux in the LHS and decreases even more when leaving this state.
+ Band-limited noise is then suppressed and quasi-periodical oscillations (QPOs) make their appearance?22?)., Band-limited noise is then suppressed and quasi-periodical oscillations (QPOs) make their appearance.
+. Not all transient BHCs in outburst go through a complete q-track., Not all transient BHCs in outburst go through a complete q-track.
+" So far, a limited number of objects were Observed to start an outburst, reach the HIMS, but then return to the LHS instead of moving to HSS The available data of the 2011 outburst of ssuggest that this event represents another example of these “failed” outbursts."," So far, a limited number of objects were observed to start an outburst, reach the HIMS, but then return to the LHS instead of moving to HSS The available data of the 2011 outburst of suggest that this event represents another example of these “failed” outbursts."
+" As summarized in Sect. 1,"," As summarized in Sect. \ref{sec:intro},"
+ the initial observations of the source carried out with MAXI indicate that the onset of the outburst occurred probably on 8802 MJD., the initial observations of the source carried out with MAXI indicate that the onset of the outburst occurred probably on 802 MJD.
+" Spectral and timing information on the source were first available through Swift//XRT and RXTE/PCA observations in the interval 8814 MJD, showing characteristics typical of the end of LHS: a relatively hard power-law photon index increasing from ~1.7 to 2.0, a relatively high RMS (42396, Fig. 3))"," Spectral and timing information on the source were first available through /XRT and RXTE/PCA observations in the interval 814 MJD, showing characteristics typical of the end of LHS: a relatively hard power-law photon index increasing from $\sim$ 1.7 to 2.0, a relatively high RMS $\sim$, Fig. \ref{fig:spectra}) ),"
+ and a marginally significant ΟΡΟ corresponding to the break frequency of the PSD., and a marginally significant QPO corresponding to the break frequency of the PSD.
+" A soft diskBB component was detected in the XRT spectrum, with a temperature ~0.3 kkeV, i..e, comparable with that expected from a truncated accretion disk in LHS(?)."," A soft diskBB component was detected in the XRT spectrum, with a temperature $\sim$ keV, i..e, comparable with that expected from a truncated accretion disk in LHS."
+". During the following bright phase of the outburst 8822 MJD), the diskBB accounts for about half the X-ray flux (Fig. 2))"," During the following bright phase of the outburst 822 MJD), the diskBB accounts for about half the X-ray flux (Fig. \ref{fig:broadband}) )"
+" with an increased temperature of ~0.4kkeV and a roughly halved squared root of normalization, which indicates a smaller inner truncation radius of the disk (Fig."," with an increased temperature of $\sim0.4$ keV and a roughly halved squared root of normalization, which indicates a smaller inner truncation radius of the disk (Fig."
+ 1 and Table 1))., \ref{fig:lcurve} and Table \ref{tab:log}) ).
+" Correspondingly, the power-law steepens (Γ~2) consistently with a more efficient cooling of a population of high-energy electrons, which up-scatter the soft disk photons and produce the high-energy non-thermal emissiontherein)."," Correspondingly, the power-law steepens $\Gamma\sim$ 2) consistently with a more efficient cooling of a population of high-energy electrons, which up-scatter the soft disk photons and produce the high-energy non-thermal emission."
+". These properties indicate a transition to the HIMS, which is confirmed by a decrease of the fractional RMS to ~18%,, the higher frequency extension of the broad-band noise, and the appearance of blandly coherent noise at a few Hertz."," These properties indicate a transition to the HIMS, which is confirmed by a decrease of the fractional RMS to $\sim$, the higher frequency extension of the broad-band noise, and the appearance of blandly coherent noise at a few Hertz."
+ The latter often evolves into QPOs with high quality factors during the HIMS?);; this is not observed forMAXIJ1836—194., The latter often evolves into QPOs with high quality factors during the HIMS; this is not observed for.
+. A transition to the SIMS is excluded by the absence of a clear drop in the RMS to a few percent (see Fig. 3)), A transition to the SIMS is excluded by the absence of a clear drop in the RMS to a few percent (see Fig. \ref{fig:spectra}) )
+ and by radio detection of the source from ~55 8806 to ~55 MMJD?)., and by radio detection of the source from $\sim$ 806 to $\sim$ MJD.
+". After 8823 MJD, uunderwent a relatively rapid flux decrease, a significant spectral hardening and an increase of the fractional RMS (see Fig. 3))."," After 823 MJD, underwent a relatively rapid flux decrease, a significant spectral hardening and an increase of the fractional RMS (see Fig.\ref{fig:spectra}) )."
+ The hard X-ray flux as measured by RXTE/PCA resumed during the first part of the decay (until ~55 8832 MJD) and then decayed at roughly constant spectral slope., The hard X-ray flux as measured by RXTE/PCA resumed during the first part of the decay (until $\sim$ 832 MJD) and then decayed at roughly constant spectral slope.
+" Correspondingly, the temperature of the thermal component decreased to ~0.2kkeV (Fig. 1))"," Correspondingly, the temperature of the thermal component decreased to $\sim$ keV (Fig. \ref{fig:lcurve}) )"
+" and its normalization increased: this can be interpreted, within the disk truncated model, as the inner radius moving away from the BH."," and its normalization increased: this can be interpreted, within the disk truncated model, as the inner radius moving away from the BH."
+ ddata collected in the period from 48 to 61 days after the onset of the event did not show evidence for the soft spectral component probably because of the limited and short exposure time of the Swift//XRT data and/or a very low disk temperature., data collected in the period from 48 to 61 days after the onset of the event did not show evidence for the soft spectral component probably because of the limited and short exposure time of the /XRT data and/or a very low disk temperature.
+" In this phase, the PL photon index remained virtually constant at L'«1.6, and the source became fainter down to an X-ray flux of 3x10! eflux.. The HID of sshowed in Fig."," In this phase, the PL photon index remained virtually constant at $\Gamma$$\simeq$ 1.6, and the source became fainter down to an X–ray flux of $\times$ $^{-10}$ The HID of showed in Fig."
+ 5 also supports these conclusions: the source, \ref{fig:hid} also supports these conclusions: the source
+Following Engineeretal.(2000) we continue to think of Rit). as defined by Eq. (1)).,"Following \citet{Engineer:2000} we continue to think of $R(t)$, as defined by Eq. \ref{deltaeqn}) ),"
+ as the effective “radius” of each shell of coustant AL., as the effective “radius” of each shell of constant $M$ .
+ The peculiar velocity. /sc. is defined by: where (=@/fa is the Hubble parameter.," The peculiar velocity, $h_{SC}$ is defined by: where $H = \dot{a}/a$ is the Hubble parameter."
+" In a matter dominated Universe the unmodified SCM predicts Pac:=Ace (8), ", In a matter dominated Universe the unmodified SCM predicts $h_{SC} = h_{SC}(\bar{\delta})$ .
+Our kev assumption in what follows is that this τομάτας the case even when deviations from spherical svinmnetry are taken into account., Our key assumption in what follows is that this remains the case even when deviations from spherical symmetry are taken into account.
+ This assumption is supported by nmunuercal studies of structure formation such as that conducted by Παποιetal.(1991) which we discuss below., This assumption is supported by numerical studies of structure formation such as that conducted by \citet{Hamilton:1991} which we discuss below.
+ For the moment we take the background universe to be matter dominated. however in Section ?? we ecucralize our results to non-lincar overdeusities in more realistic nuiverses Which have receutly trausitioned to an epoch of dark cuerey domination.," For the moment we take the background universe to be matter dominated, however in Section \ref{degen} we generalize our results to non-linear overdensities in more realistic universes which have recently transitioned to an epoch of dark energy domination."
+ We define a new coordinate (0) by: for some function 147)., We define a new coordinate $\eta(\bar{\delta})$ by: for some function $T(\eta)$ .
+ Now M is constant on cach shell aud so. taking f£.= f;. R=R; and 6=39; at some instance. we lave By considering this relation together with Eq. (3))," Now $M$ is constant on each shell and so, taking $t=t_{i}$ , $R=R_{i}$ and $\bar{\delta}=\delta_{i}$ at some instance, we have By considering this relation together with Eq. \ref{deltaeqnT}) )"
+ we see that: where j(4.0;) is some function of aud à;.," we see that: where $j(\eta, \delta_i)$ is some function of $\eta$ and $\delta_i$."
+ If hac:= then jGpRi)=jp. aud without loss of generality we can set j(4)=1.," If $h_{SC} = h_{SC}(\bar{\delta})$ then $j(\eta, R_i)=j(\eta)$, and without loss of generality we can set $j(\eta) = 1$."
+ This gives: where το=4g—sinj. aud: We reduce to the standard SCALD in the luit where T—r(g)-nmsing (Padmanabhan2002).. and it is clear that ορ.=ge(6) in the mumocdified SCAL," This gives: where $\tau(\eta) = \eta - \sin \eta$, and: We reduce to the standard SCM in the limit where $T = \tau(\eta) =\eta - \sin \eta$ \citep{Pad:2002}, and it is clear that $h_{SC} = h_{SC}(\bar{\delta})$ in the unmodified SCM."
+ It is iurportant to stress that if Ages=Ave(6) then Eqs. (5)), It is important to stress that if $h_{SC} = h_{SC}(\bar{\delta})$ then Eqs. \ref{Reta}) )
+ and (61) are the most ecucral solutions for &@ aud f., and \ref{Teta}) ) are the most general solutions for $R$ and $t$.
+" The radial acceleration of each shell is fouud to be: where 1)=dT/dz. and T""=qT/dz.As should be expected. when T!=1 we reduce to the standard SCAL equation for Ag."," The radial acceleration of each shell is found to be: where $T^{\prime} = \dd T / \dd \tau$, and $T^{\prime \prime} = \dd^2 T / \dd\tau^2$.As should be expected, when $T^{\prime} = 1$ we reduce to the standard SCM equation for $R_{,tt}$."
+ When deviations from spherical sviuuetry and the leading order effects of a gradient iu the velocity dispersion are taken iuto account. Eneiuecretal.(2000) showed that one should have: where οί).=«tla?207j|μυ c? aud QO? respectively quautifv the shear and rotation of the fluid: f(o..0) contains thelowest order coutribution frou velocity dispersion terms.," When deviations from spherical symmetry and the leading order effects of a gradient in the velocity dispersion are taken into account, \citet{Engineer:2000} showed that one should have: where $S(a, x) = a^2(\sigma^2 - 2\Omega^2) + f(a,x)$ ; $\sigma^2$ and $\Omega^2$ respectively quantify the shear and rotation of the fluid; $f(a,x)$ contains thelowest order contribution from velocity dispersion terms."
+ huportautlv. no matter what form S(a..c) takes. if as we have assumed σοι=hse). then we must have 5=$(d) and by comparing Eqs. Cr))," Importantly, no matter what form $S(a,x)$ takes, if, as we have assumed $h_{SC}= h_{SC}(\bar{\delta})$, then we must have $S = S(\bar{\delta})$ and by comparing Eqs. \ref{modeqn}) )"
+ aud (8)) we can clearly see that: Tn principle. the form of both $(6) aud hence Ty) can be found using the results of N-body simulations.," and \ref{fluideqn}) ) we can clearly see that: In principle, the form of both $S(\bar{\delta})$ and hence $T(\eta)$ can be found using the results of N-body simulations."
+ Uufortunatev. however. making the required comparison with simulations is uot as straightforward as one nüeht expect it to be.," Unfortunately, however, making the required comparison with simulations is not as straightforward as one might expect it to be."
+ This is because the results of such simulations are given iu terms of the statistical properties of the matter distribution rather than iu terms of the mean density contrast. ὃν aud peculiar velocity. lise.," This is because the results of such simulations are given in terms of the statistical properties of the matter distribution rather than in terms of the mean density contrast, $\bar{\delta}$, and peculiar velocity, $h_{SC}$."
+ The statistical properties in question are the averaged two point correlation function. £. and the averaged pair velocity. h(a..c). which are given as defined by: where © is the two-poimt correlation function aud isdefiued to be the Fourier tfransforii of the power spectrum. P(A}.," The statistical properties in question are the averaged two point correlation function, $\bar{\xi}$, and the averaged pair velocity, $h(a,x)$ , which are given as defined by: where $\xi$ is the two-point correlation function and isdefined to be the Fourier transform of the power spectrum, $P(k)$."
+ The assumption that Ata.) depends ou « ande onlv through & Lie. Ata.)=hb(£(a..0)) is ΟΠΛΟ in the literature (see Engineeretal.(2000).. Moeetal. (1995))) aud it appears to have been confined by nunercal simulations (see Ibuuitouetal.(1991).. Peacock&Dodds (1996))," The assumption that $h(a,x)$ depends on $a$ and $x$ only through $\bar{\xi}$ i.e. $h(a,x) = h(\bar{\xi}(a,x))$ is common in the literature (see \cite{Engineer:2000}, \cite{Moe:1995}) ) and it appears to have been confirmed by numerical simulations (see \cite{Hamilton:1991}, \cite{Peacock:1996}) )."
+ The results of. for example.). Hamiltonctal.(1991) can be used to construct the fitting formula for h(£) (Eneinecretal.2000).," The results of, for example, \citet{Hamilton:1991} can be used to construct the fitting formula for $h(\bar{\xi})$ \citep{Engineer:2000}."
+. Towever. before we can make use of such a formula. we iust relate the statistical quautities & aud h(£) to 6 aud hae(0).," However, before we can make use of such a formula, we must relate the statistical quantities $\bar{\xi}$ and $h(\bar{\xi})$ to $\bar{\delta}$ and $h_{SC}(\bar{\delta})$ ."
+ It is well known that (Padmanabhan (1996)..Pacinauabhau (2002)..Peebles (1980)..Pacinanabhan&Eneinecr(1996)..Popoloetal. (2001)..Padmanabhan&Rav (2006))). 01 scalessmaller than the size of the collapsing objects and around biel density peaks: It follows that. in the non-linear regime. we have 6zx &.," It is well known that \cite{Pad:1996}, \cite{Pad:2002}, \cite{Peebles:1980}, \cite{Padeng},\cite{pop}, \cite{pad}) ), on scalessmaller than the size of the collapsing objects and around high density peaks: It follows that, in the non-linear regime, we have $\bar{\delta} \approx \bar{\xi}$ ."
+ This relationship between 6 and © was also used by Eneiecretal. (2000).. although it was. as it is here. the weakest partof the whole analysis.," This relationship between $\bar{\delta}$ and $\bar{\xi}$ was also used by \citet{Engineer:2000}, , although it was, as it is here, the weakest partof the whole analysis."
+ Weouly require that&z© hold where it is expected tobe a good approximation ic. 6z 15., Weonly require that$\bar{\delta} \approx \bar{\xi}$ hold where it is expected tobe a good approximation i.e. $\bar{\delta} \gtrsim 15$ .
+ As 6>»o we assume that à~ & Peebles(1980). showed that 52)satisfies:, As $\bar{\delta} \rightarrow \infty$ we assume that $\bar{\delta} \sim \bar{\xi}$ \citet{Peebles:1980} showed that $h(\bar{\xi})$satisfies:
+lines (visible even. at SAIC metallicity): this behaviour is most obvious in the ratio of A to Lis. but is clear in other features. such as the A44300 C-band.,"lines (visible even at SMC metallicity); this behaviour is most obvious in the ratio of $\;K$ to $\gamma$, but is clear in other features, such as the 4300 $G$ -band."
+ The adopted Jate-type classification. scheme. which takes into account the characteristics of the 201 dataset (low sin. weak metal lines). is summoarised in Table 4," The adopted `late-type' classification scheme, which takes into account the characteristics of the 2dF dataset (low s:n, weak metal lines), is summarised in Table \ref{fclass}."
+ For population svntheses. LAL modelling. and many other circumstances where quantitative analyses olf individual stars is infeasible. it is important to have an accurate ‘look-up table’ of correspondences between physical. and morphological characteristics.," For population syntheses, IMF modelling, and many other circumstances where quantitative analyses of individual stars is infeasible, it is important to have an accurate `look-up table' of correspondences between physical and morphological characteristics."
+ 1n the remainder of this paper. we examine this issue in the light of our remarks in section 2.. with particular emphasis on the temperature scale for A-type stars.," In the remainder of this paper, we examine this issue in the light of our remarks in section \ref{principles}, with particular emphasis on the temperature scale for A-type stars."
+ The model-atmosphere structures used. here are models (?).. ealeulated by Collaborative Computational project 77 (CCRT: http://ccp7.dur.ac.uk/).," The model-atmosphere structures used here are models \citep{k91}, calculated by Collaborative Computational Project 7 (CCP7; )."
+ The models assume plane-parallel geometry. hyclrostatic equilibrium.— anc LEE. and include the elfects.— of line danketing.," The models assume plane-parallel geometry, hydrostatic equilibrium, and LTE, and include the effects of line blanketing."
+ “Phe majority of the A-type stars in the 2dE sample are of luminosity tvpes of Ib or H. where LTIS has oen shown to be a reasonable approximation (c.g.?)..," The majority of the A-type stars in the 2dF sample are of luminosity types of Ib or II, where LTE has been shown to be a reasonable approximation \citep[e.g.][]{pryz02}."
+ Departures from L'TIZ. become increasingly important in ype la supergiants: however. in the current. context. theΑΟ models have the advantage that results can be achieved quickly. relatively simply. anc that they include he important effects of line blanketing.," Departures from LTE become increasingly important in type Ia supergiants; however, in the current context, the models have the advantage that results can be achieved quickly, relatively simply, and that they include the important effects of line blanketing."
+ The. low-metallicity CCRT models that we κου incorporate a microturbulent velocity. £. of 2 +.," The low-metallicity CCP7 models that we used incorporate a microturbulent velocity, $\xi$, of 2 ."
+" ""This is just below the range of€ = 3S found in Venn’s A-supergiant analyses (7.1999).. although many of her targets vielded. values in the 34 range."," This is just below the range of $\xi$ = 3–8 found in Venn's A-supergiant analyses \citep[][1999]{venn95}, although many of her targets yielded values in the 3–4 range."
+ Spectra svnthesized using 2- and|. solar-abuncdance models are essentially identical. ane cdillerences resulting from use of an &erid are negligible.," Spectra synthesized using 2- and, solar-abundance models are essentially identical, and differences resulting from use of an grid are negligible."
+ We conclude that the CCP models are adequate for the current investigation., We conclude that the CCP7 models are adequate for the current investigation.
+ Llubeny’s unpublished. code was used for the spectral synthesis. running under the wrapper (also written by Hubeny).," Hubeny's unpublished code was used for the spectral synthesis, running under the wrapper (also written by Hubeny)."
+ downloadedThe necessary atomic and molecular line-lists were from the web Both Ravleigh-seattering ancl -ion opacities were included in the calculations., The necessary atomic and molecular line-lists were downloaded from the web Both Rayleigh-scattering and $^-$ -ion opacities were included in the calculations.
+ The synthetic spectra were convolved. with typical values (~30 *)). and smoothed to a resolution comparable to that of the 01) data.," The synthetic spectra were convolved with typical values $\sim{30}$ ), and smoothed to a resolution comparable to that of the 2dF data."
+ In all cases. instrumental broadening dominates.," In all cases, instrumental broadening dominates."
+ Three main spectral regions were investigated: The atomic species included. for the spectral-svnthesis investigation of B and A-twpes were Η. He. €. Ν. O. Mg. οἱ and Ca.," Three main spectral regions were investigated: The atomic species included for the spectral-synthesis investigation of B and A-types were H, He, C, N, O, Mg, Si and Ca."
+ For the later tvpes. in addition to the molecular line-list. An. Cr and Fe were included.," For the later types, in addition to the molecular line-list, Mn, Cr and Fe were included."
+ was used to generate svnthetie spectra on a small erid covering appropriate values of temperature and eravitv., was used to generate synthetic spectra on a small grid covering appropriate values of temperature and gravity.
+ The temperature range of the eric was chosen by reference to ?.., The temperature range of the grid was chosen by reference to \citet{sk82}.
+ Onlv approximate values are required for the sampling because temperature has a more significant cllect on the lines studied here., Only approximate values are required for the sampling because temperature has a more significant effect on the lines studied here.
+" Typical values of for A-type supergiants were estimated from ?. and οι, ?.", Typical values of for A-type supergiants were estimated from \citet{venn95} and \citet{vtg99b}.
+ derive 8 2.3 for two BS lab stars. so = 2.5 was taken as an upper limit for the hotter models.," \citet{duf00}
+ derive $\approx~$ 2.3 for two B5 Iab stars, so = 2.5 was taken as an upper limit for the hotter models."
+"g@ For the E ancl CG spectral types considered here. ?. find values over a range ollog,,g = 0.01.0.so spectra were svnthesized. for our Coolest models using = "," For the F and G spectral types considered here, \citet{lmb97} find values over a range of = 0.0–1.0,so spectra were synthesized for our coolest models using = 0.5."
+The grid samplingis shown in Figure 6.., The grid samplingis shown in Figure \ref{grid}. .
+ bor our calculations we have adopted à species-mndependent, For our calculations we have adopted a species-independent
+Froii the discussion im the precediug sections. it follows that equation (20)) holds when the adiabatic iudex is very close to. 1/3 aud οελαX10.,"From the discussion in the preceding sections, it follows that equation \ref{eqn-fields}) ) holds when the adiabatic index is very close to 4/3 and $M(>\epsilon)/M_{\rm ej}\ltsim 10^{-5}$."
+ Iu particular. our caleulatious show that the enerev distribution docs not become a simple power-law for M(26)/M4Z10.," In particular, our calculations show that the energy distribution does not become a simple power-law for $M(>\epsilon)/M_{\rm ej}\gtsim 10^{-5}$."
+ This must be due to the fact that the simplestellar density distribution Matzuer&Melee(1999) assiunced can reproduce the deusity distributions iu realistic stellar model only for the outer 0.001 lavers in mass., This must be due to the fact that the simplestellar density distribution \citet{Matzner_99} assumed can reproduce the density distributions in realistic stellar model only for the outer 0.001 layers in mass.
+ When the explosion energy is furious. the adiabatic iudex is close to 1/3 but the mass with e ereater than the threshold euergies for spallation reactions can become significantly ereatcr than 10M as iu SN L998bw model.," When the explosion energy is furious, the adiabatic index is close to 4/3 but the mass with $\epsilon$ greater than the threshold energies for spallation reactions can become significantly greater than $10^{-5}M_{\rm ej}$ as in SN 1998bw model."
+ Thus equation (203) does not eive a good estimate for the mass of ejecta iu the region between e aud €|de around the threshold euergies dM e\/de) as well as Mt €)., Thus equation \ref{eqn-fields}) ) does not give a good estimate for the mass of ejecta in the region between $\epsilon$ and $\epsilon+d\epsilon$ around the threshold energies $-dM(>\epsilon)/d\epsilon$ ) as well as $M(>\epsilon)$ .
+"A typical choice for W is a circular Gaussian centered at the galactic centroid, Alternatively, one can choose to optimize the weight function to the shape of the source to be measured.","A typical choice for $W$ is a circular Gaussian centered at the galactic centroid, Alternatively, one can choose to optimize the weight function to the shape of the source to be measured."
+" ?,seetheirsection proposed the usage of a Gaussian, whose centroid x,, size s, and ellipticity € are matched to the source, such that the argument of the exponential in is modified according to As such a weight function represents a matched spatial filter, it optimizes the significance and accuracy of the measurement if its parameters are close to their true values."," \citet[see their section 3.1.2]{Bernstein02.1} proposed the usage of a Gaussian, whose centroid $\mathbf{x}_c$, size $s$, and ellipticity $\bepsilon$ are matched to the source, such that the argument of the exponential in is modified according to As such a weight function represents a matched spatial filter, it optimizes the significance and accuracy of the measurement if its parameters are close to their true values."
+" This can, however, not be guaranteed in presence of pixel noise, but we found the iterative algorithm proposed by ? to converge well in practice and therefore employ it to set the weight function within the method."," This can, however, not be guaranteed in presence of pixel noise, but we found the iterative algorithm proposed by \citet{Bernstein02.1} to converge well in practice and therefore employ it to set the weight function within the method."
+" Unfortunately, a product in real space like the one in translates into a convolution in Fourier-space."," Unfortunately, a product in real space like the one in translates into a convolution in Fourier-space."
+" We therefore have to expect some amount of mixing of the moments of /,,.", We therefore have to expect some amount of mixing of the moments of $I_w$.
+" Even worse, an attempt to relate the moments of /,, to those of J leads to diverging integrals."," Even worse, an attempt to relate the moments of $I_w$ to those of $I$ leads to diverging integrals."
+" Hence, there is no exact way of incorporating spatial weighting to the moment approach outlined above."," Hence, there is no exact way of incorporating spatial weighting to the moment approach outlined above."
+" On the other hand, we can invert for J=J,,/W and expand 1/W ina Taylor series around the center at x—0, where we employed W’(x)= andciΞ(1τει)’+e."," On the other hand, we can invert for $I=I_w/W$ and expand $1/W$ ina Taylor series around the center at $\mathbf{x}=\mathbf{0}$, where we employed $W'(\mathbf{x})\equiv \frac{dW(\mathbf{x})}{d\mathbf{x}^2}$ and $c_{1,2} \equiv (1\mp\epsilon_1)^2 + \epsilon_2^2$."
+" We introduce the parameter n,, as the aemaximum order of the Taylor expansion, here n,= 4."," We introduce the parameter $n_w$ as the maximum order of the Taylor expansion, here $n_w=4$ ."
+" Inserting this expansion in1,, we are able to approximate the moments of / by their counterparts {/,,,}."," Inserting this expansion in, we are able to approximate the moments of $I$ by their counterparts $\mom{I_{dw}}{}$."
+" For convenience we give the correction terms for orders n,,<6 in2.", For convenience we give the correction terms for orders $n_w \leq 6$ in.
+. This linear expansion allows us to correct for the weighting-induced change in the moments of a certain order n by considering the impact of the weight function on weighted moments up to order n+ny., This linear expansion allows us to correct for the weighting-induced change in the moments of a certain order $n$ by considering the impact of the weight function on weighted moments up to order $n+n_w$.
+ The truncation of the Taylor expansion constitutes the first and only source of bias in the method., The truncation of the Taylor expansion constitutes the first and only source of bias in the method.
+" The direction of the bias is evident: As the weight function suppresses contributions to the moments from pixel far away from the centroid, its employment reduces the power in any moment by an amount, which depends on the shape — particularly on the radial profile — of the source and the width s."," The direction of the bias is evident: As the weight function suppresses contributions to the moments from pixel far away from the centroid, its employment reduces the power in any moment by an amount, which depends on the shape – particularly on the radial profile – of the source and the width $s$."
+" Additionally, if the ellipticity e was misestimated during the matching of W, the measured ellipticity of the source x before and after deweighting will be biased towards e."," Additionally, if the ellipticity $\bepsilon$ was misestimated during the matching of $W$, the measured ellipticity of the source $\bchi$ before and after deweighting will be biased towards $\bepsilon$."
+" In the realistic case of noisy images, which we address in more detail in3.2,, e can be wrong in two ways: statistically and systematically."," In the realistic case of noisy images, which we address in more detail in, $\bepsilon$ can be wrong in two ways: statistically and systematically."
+" The centered Gaussian distribution of the pixel noise leads to a centered Cauchy-type error distribution of both components of e, i.e. the statistical errors of € and therefore also y have vanishing mean."," The centered Gaussian distribution of the pixel noise leads to a centered Cauchy-type error distribution of both components of $\bepsilon$, i.e. the statistical errors of $\bepsilon$ and therefore also $\chi$ have vanishing mean."
+" The systematic errors stem from the application of a compact weight function to measure €, which for any finite s constitutes a removal of information."," The systematic errors stem from the application of a compact weight function to measure $\epsilon$, which for any finite $s$ constitutes a removal of information."
+ This leads to deviations of the measured e from the true e if e.g. the centroid is not determined accurately or the ellipticity changes with radius., This leads to deviations of the measured $\bepsilon$ from the true $\bepsilon$ if e.g. the centroid is not determined accurately or the ellipticity changes with radius.
+" For individual objects, these deviations are impossible to quantify precisely — as this would require knowledge of the true observed morphology — and hamper all weak-lensing measurements unless an appropriate treatment is devised (e.g.??).."," For individual objects, these deviations are impossible to quantify precisely – as this would require knowledge of the true observed morphology – and hamper all weak-lensing measurements unless an appropriate treatment is devised \citep[e.g.][]{Hosseini09.1, Bernstein10.1}."
+" For large ensembles of galaxies, effective levels of systematic errors can be assessed by analyzing dedicated simulations (cf. 4))."," For large ensembles of galaxies, effective levels of systematic errors can be assessed by analyzing dedicated simulations (cf. )."
+" We investigate now the -specific systematic impact of a finite n, on the recovery of the deweighted moments.", We investigate now the -specific systematic impact of a finite $n_w$ on the recovery of the deweighted moments.
+" For the experiments inthis section we simulated simple galaxy models following the Sérrsic profile where n, denotes the Sérrsic index, and R, the effective radius, while the PSFs are modeled from the Moffat profile where a=(2!/8—1)/(FWHM/2) sets the width of the profile and f its slope."," For the experiments inthis section we simulated simple galaxy models following the Sérrsic profile where $n_s$ denotes the Sérrsic index, and $R_e$ the effective radius, while the PSFs are modeled from the Moffat profile where $\alpha =
+\bigl(2^{1/\beta}-1\bigl)/\bigl(\mathrm{FWHM}/2\bigr)^2$ sets the width of the profile and $\beta$ its slope."
+ Both model types acquire their ellipticity according to In the top panel of we show the error after deweighting a convolved galaxy image from a matched elliptical weight function as a function of its size s., Both model types acquire their ellipticity according to In the top panel of we show the error after deweighting a convolved galaxy image from a matched elliptical weight function as a function of its size $s$.
+" As noted above, the bias is always negative and is clearly more prominent for the larger disk-type galaxy (circle markers)."," As noted above, the bias is always negative and is clearly more prominent for the larger disk-type galaxy (circle markers)."
+" As the Taylor expansion becomes more accurate for n,—oo or 5—oo, the bias of any moment decreases accordingly."," As the Taylor expansion becomes more accurate for $n_w\to\infty$ or $s\to\infty$, the bias of any moment decreases accordingly."
+" An important consequence of the employment of a weight function with matched ellipticity is that the bias after deweighting does only very weakly depend on the apparent ellipticity, i.e. all moments of the same order are biased by the same relative factor A(n,s)."," An important consequence of the employment of a weight function with matched ellipticity is that the bias after deweighting does only very weakly depend on the apparent ellipticity, i.e. all moments of the same order are biased by the same relative factor $\Delta(n,s)$."
+ This means any ratio of such moments remains unbiased., This means any ratio of such moments remains unbiased.
+" This does not guarantee that the ellipticity is still unbiased after the moments have passed the deconvolution step, which is exact only for unweighted moments."," This does not guarantee that the ellipticity is still unbiased after the moments have passed the deconvolution step, which is exact only for unweighted moments."
+" On the other hand, the particular form of the equations in becomes important here: If we assume well-centered images of the galaxy and the PSF and negligible error of the source flux (G]oo, none of which isguaranteeda for faint objects, the deconvolution equations for the relevant second- moments only mix second-order moments."," On the other hand, the particular form of the equations in becomes important here: If we assume well-centered images of the galaxy and the PSF and a negligible error of the source flux $\mom{G}{0,0}$ , none of which isguaranteed for faint objects, the deconvolution equations for the relevant second-order moments only mix second-order moments."
+" If furthermore Δα(2.s)=Ap(2, s), the ellipticity x (cf. 2))"," If furthermore $\Delta_G(2,s)=\Delta_P(2,s)$ , the ellipticity $\bchi$ (cf. )"
+ will remain unbiased after deconvolution even though the moments themselves, will remain unbiased after deconvolution even though the moments themselves
+test.,test.
+" For the semirelativistic model with fixed ""y for the subsample with Rmaz=10000kms~' the minimum value F=11.5 corresponds to Ci, the maximum value F=109.7 — to Co."," For the semirelativistic model with fixed $\gamma$ for the subsample with $R_{max}=10000\,\rmn{km\,s}^{-1}$ the minimum value $F=11.5$ corresponds to $C_1$ , the maximum value $F=109.7$ – to $C_2$."
+" These values should be compared to the values 3.8, 6.6, 7.9, 10.8 and 12.1, which correspond to 95, 99, 99.5, 99.9 and 99.95 per cent confidence levels respectively."," These values should be compared to the values $3.8$, $6.6$ , $7.9$ , $10.8$ and $12.1$, which correspond to $95$, $99$, $99.5$, $99.9$ and $99.95$ per cent confidence levels respectively."
+" Thus, all the coefficients of the generalised Tully-Fisher relation (16)) are statistically significant at the 99.9 per cent confidence level."," Thus, all the coefficients of the generalised Tully-Fisher relation \ref{eqn:TFR}) ) are statistically significant at the $99.9$ per cent confidence level."
+ Now let us consider the dipole component of the velocity field., Now let us consider the dipole component of the velocity field.
+" Its parameters including the galactic coordinates /,b of the apex for the DQO-model are given in Table 1.."," Its parameters including the galactic coordinates $l, b$ of the apex for the DQO-model are given in Table \ref{tbl:1}."
+ The norms of the dipolar component do not contradict the ACDM model., The norms of the dipolar component do not contradict the $\Lambda$ CDM model.
+" For the model with fixed ; the module of the dipolar component drops to 180kms""."," For the model with fixed $\gamma$ the module of the dipolar component drops to $180\,\rmn{km\,s}^{-1}$."
+" However, the bulk motion is usually considered in the framework of the simplest dipole models when the only characteristics of the velocity field are the modulus and the apex of the dipole component."," However, the bulk motion is usually considered in the framework of the simplest dipole models when the only characteristics of the velocity field are the modulus and the apex of the dipole component."
+ In our case of DQO-models the velocity field is more complex and we cannot attribute the bulk motion solely to the dipole component., In our case of DQO-models the velocity field is more complex and we cannot attribute the bulk motion solely to the dipole component.
+" For this reason, to compare ourresults to the results of other authors we also calculated the dipole component in the framework of the relativistic (the same as semirelativistic) model with fixed y."," For this reason, to compare ourresults to the results of other authors we also calculated the dipole component in the framework of the relativistic (the same as semirelativistic) D-model with fixed $\gamma$ ."
+" It yields the bulk flow velocity of 314kms! directed towards |=322?,b27° (Centaurus)."," It yields the bulk flow velocity of $314\,\rmn{km\,s}^{-1}$ directed towards $l=322\degr, b=27\degr$ (Centaurus)."
+" On Figure 3 we plotted the boundaries of 1c, 2c and 3c confidence areas of this apex for Rmax=10000kms."," On Figure \ref{fig:3} we plotted the boundaries of $1\sigma$, $2\sigma$ and $3\sigma$ confidence areas of this apex for $R_{max}=10000\,\rmn{km\,s}^{-1}$."
+ For this purpose we projected the 8-dimensional ellipsoid of errors into the 3-dimensional space and then projected it on the celestial sphere., For this purpose we projected the 8-dimensional ellipsoid of errors into the 3-dimensional space and then projected it on the celestial sphere.
+ On the same figure we also plotted the boundaries of the confidence areas of the apex in non-relativistic D-model (Parnovsky&Parnowski2010) as well as positions of apices obtained by different authors., On the same figure we also plotted the boundaries of the confidence areas of the apex in non-relativistic D-model \citep{ref:APSS09} as well as positions of apices obtained by different authors.
+ The value of the bulk motion appears to be larger than for DQO-models., The value of the bulk motion appears to be larger than for DQO-models.
+" For the subsample with Rinaz=8000kms'! it is equal to 285kms""!."," For the subsample with $R_{max}=8000\,\rmn{km\,s}^{-1}$ it is equal to $285\,\rmn{km\,s}^{-1}$."
+" We see that D-models provide a result, which is closer to that obtained by Watkinsetal.(2009), but still consistent with the ACDM model."," We see that D-models provide a result, which is closer to that obtained by \citet{ref:WFH09}, but still consistent with the $\Lambda$ CDM model."
+ Let us now consider the quadrupole component of the velocity field., Let us now consider the quadrupole component of the velocity field.
+ What is the physical sense of the quadrupole component?, What is the physical sense of the quadrupole component?
+" As one can see from the paper (Parnovskyetal.2001),, it can be naturally combined with the Hubble constant."," As one can see from the paper \citep{ref:Par01}, it can be naturally combined with the Hubble constant."
+" As a result, we obtain the effective ‘Hubble constant’ depending on direction Naturally, this effective ‘Hubble constant’ is caused by the large-scale collective motion on the sample scale."," As a result, we obtain the effective `Hubble constant' depending on direction Naturally, this effective `Hubble constant' is caused by the large-scale collective motion on the sample scale."
+ To estimate the value of its anisotropy we found the eigenvalues and eigenvectors of tensor Q., To estimate the value of its anisotropy we found the eigenvalues and eigenvectors of tensor $Q$.
+ The three eigenvectors are orthogonal and the sum of three eigenvalues is equal to zero because Q is a traceless tensor., The three eigenvectors are orthogonal and the sum of three eigenvalues is equal to zero because $\tens{Q}$ is a traceless tensor.
+ We found the eigenvalues and the eigenvectors of the tensor Q for two considered subsamples., We found the eigenvalues and the eigenvectors of the tensor $\tens{Q}$ for two considered subsamples.
+" For the subsample with Rmax=10000kms! the maximal eigenvalue 7.1+1.7 per cent corresponds to an axis directed towards |=98°,b79° (Canes Venatici) and the opposite direction (Phoenix)."," For the subsample with $R_{max}=10000\,\rmn{km\,s}^{-1}$ the maximal eigenvalue $7.1 \pm 1.7$ per cent corresponds to an axis directed towards $l=98\degr, b=79\degr$ (Canes Venatici) and the opposite direction (Phoenix)."
+" The minimal eigenvalue —4.1+1.7 per cent corresponds to an axis directed towards 1=195°,b2° (Gemini) and the opposite direction (Sagittarius)."," The minimal eigenvalue $-4.1 \pm 1.7$ per cent corresponds to an axis directed towards $l=195\degr, b=2\degr$ (Gemini) and the opposite direction (Sagittarius)."
+" The third eigenvalue —3 per cent corresponds to an axis directed towards |=286°,b11? (Centaurus-Vela) and the opposite direction (Andromeda-Lacerta)."," The third eigenvalue $-3$ per cent corresponds to an axis directed towards $l=286\degr, b=11\degr$ (Centaurus-Vela) and the opposite direction (Andromeda-Lacerta)."
+ Comparing these values to the non-relativistic model (Parnovsky&Parnowski2010) one can see that both the eigenvalues and the directions of the axes changed insignificantly., Comparing these values to the non-relativistic model \citep{ref:APSS09} one can see that both the eigenvalues and the directions of the axes changed insignificantly.
+" Nevertheless, the two negative eigenvalues, which are close to each other, have the opposite order in these two models."," Nevertheless, the two negative eigenvalues, which are close to each other, have the opposite order in these two models."
+" In this sense, the positive axis notably stands out, for which the effective ‘Hubble constant’ exceeds the mean value by 7 per cent."," In this sense, the positive axis notably stands out, for which the effective `Hubble constant' exceeds the mean value by 7 per cent."
+ For the subsample with Rmaz=8000kmst the ellipsoid is three-axial and essentially differs from the oblate spheroid.," For the subsample with $R_{max}=8000\,\rmn{km\,s}^{-1}$ the ellipsoid is three-axial and essentially differs from the oblate spheroid."
+" The maximal eigenvalue 7.1+1.7 per cent corresponds to an axis directed towards |=102°,b85° (Canes Venatici) and the opposite direction (Sculptor)."," The maximal eigenvalue $7.1 \pm
+1.7$ per cent corresponds to an axis directed towards $l=102\degr, b=85\degr$ (Canes Venatici) and the opposite direction (Sculptor)."
+" The minimal eigenvalue —5.3+1.6 per cent corresponds to an axis directed towards |=266°,b5° (Vela) and the opposite direction (Cygnus)."," The minimal eigenvalue $-5.3 \pm 1.6$ per cent corresponds to an axis directed towards $l=266\degr,
+b=5\degr$ (Vela) and the opposite direction (Cygnus)."
+" The third eigenvalue —2 per cent corresponds to an axis directed towards |=356°,b1° (Sagittarius-Scorpio) and the opposite direction (Auriga-Taurus)."," The third eigenvalue $-2$ per cent corresponds to an axis directed towards $l=356\degr, b=1\degr$ (Sagittarius-Scorpio) and the opposite direction (Auriga-Taurus)."
+ These values are very close to those given by the non-relativistic model., These values are very close to those given by the non-relativistic model.
+ Note that the axes for both subsamples nearly coincide with the exception of reverse order of negative eigenvalues for Rmaz=10000kmsvt.," Note that the axes for both subsamples nearly coincide with the exception of reverse order of negative eigenvalues for $R_{max}=10000\,\rmn{km\,s}^{-1}$."
+ We also calculated the statistical significance of these eigenvalues., We also calculated the statistical significance of these eigenvalues.
+" For Rmax=10000kms! the maximal eigenvalue has F=18.2, which means that it is non-zero at 99.95 per cent confidence level, and the minimal eigenvalue has F=5.8, which means that it is non-zero at 97.5 per cent confidence level."," For $R_{max}=10000\,\rmn{km\,s}^{-1}$ the maximal eigenvalue has $F=18.2$, which means that it is non-zero at $99.95$ per cent confidence level, and the minimal eigenvalue has $F=5.8$, which means that it is non-zero at $97.5$ per cent confidence level."
+ The similar situation holds for the subsample with Rmaz=8000kms! with Fisher values being 18.3 and 10.7 respectively.," The similar situation holds for the subsample with $R_{max}=8000\,\rmn{km\,s}^{-1}$ with Fisher values being $18.3$ and $10.7$ respectively."
+" Additionally, we calculated the total statistical significance of the quadrupole component."," Additionally, we calculated the total statistical significance of the quadrupole component."
+" The value V?"" with its 5 degrees of freedom appears to be non-zero at over 99.5 per cent confidence level according to F-test.", The value $V^{qua}$ with its 5 degrees of freedom appears to be non-zero at over $99.5$ per cent confidence level according to F-test.
+ In the same way we calculated the total statistical significance of the octopole component., In the same way we calculated the total statistical significance of the octopole component.
+ The value V?** with its 10 degrees of freedom appears to be non-zero at over 99.5 per cent confidence level according to F-test., The value $V^{oct}$ with its 10 degrees of freedom appears to be non-zero at over $99.5$ per cent confidence level according to F-test.
+ The value P with its 3 degrees offreedom appears to be non-zero at slightly lessthan 90per cent confidence level according to F-test., The value $\vec{P}$ with its 3 degrees offreedom appears to be non-zero at slightly lessthan $90$per cent confidence level according to F-test.
+" Unlike the quadrupole component, the octopoleone lacks easily interpretable characteristics like eigenvector apices."," Unlike the quadrupole component, the octopoleone lacks easily interpretable characteristics like eigenvector apices."
+" The radial velocity field for R=8000kms! and Rmaz=10000kms~! in the semirelativistic model, which includes the octopole component, appeared to be very similar to that in non-relativistic case, depicted on Fig."," The radial velocity field for $R=8000\,\rmn{km\,s}^{-1}$ and $R_{max}=10000\,\rmn{km\,s}^{-1}$ in the semirelativistic model, which includes the octopole component, appeared to be very similar to that in non-relativistic case, depicted on Fig."
+ 6 in the article (Parnovsky&Parnowski 2010).., 6 in the article \citep{ref:APSS09}. .
+ The most prominent feature of, The most prominent feature of
+"Galaxy Redshift Survey (2dFGRS;Collessetal.2001) and the first-year Wilkinson Microwave Anisotropy Probe data (WMAP;Spergeletal. 2003):: Oy= 0.25, Qa= 0.75, h=0.73, n=1, and og=0.9.","Galaxy Redshift Survey \citep[2dFGRS;][]{Colless01} and the first-year Wilkinson Microwave Anisotropy Probe data \citep[WMAP;][]{Spergel03}: : $\Omega_0=0.25$ , $\Omega_\Lambda=0.75$ , $h=0.73$, $n=1$, and $\sigma_8=0.9$."
+ The haloes are found by a two-step procedure., The haloes are found by a two-step procedure.
+" First, all collapsed haloes with at least 20 particles are identified using a standard friends-of-friends group-finder with linking parameter b=0.2."," First, all collapsed haloes with at least 20 particles are identified using a standard friends-of-friends group-finder with linking parameter $b=0.2$."
+" Then, post-processing with the substructure algorithm (Springeletal.2001) subdivides each friends-of-friends halo into a set of self-bound subhaloes."," Then, post-processing with the substructure algorithm \citep{Springel01} subdivides each friends-of-friends halo into a set of self-bound subhaloes."
+ We note that comparable halo properties are found using other structure finders (seeKnebeetal. 2011)., We note that comparable halo properties are found using other structure finders \citep[see][]{Knebe11}.
+". From the Millennium Simulation halo merger tree at z= 0, we construct a mock galaxy catalogue using the halo occupation method described in Skibbaetal.(2006,here-afterS06) and Skibba&Sheth(2009,hereafter SS09);; we refer the reader to these papers for details."," From the Millennium Simulation halo merger tree at $z=0$ , we construct a mock galaxy catalogue using the halo occupation method described in \citet[][ hereafter S06]{Skibba06} and \citet[][ hereafter SS09]{Skibba09b}; ; we refer the reader to these papers for details."
+" Other halo-model descriptions of galaxy clustering—conditional luminosity functions (e.g.Yang,Mo&vandenBosch2003) and subhalo abundance matching (e.g.Kravtsovetal.2004)- —would produce similar mock catalogues, although an advantage of the SS09 approach is that it includes a strongly constrained model of galaxy colours."," Other halo-model descriptions of galaxy clustering—conditional luminosity functions \citep[e.g.][]{Yang03} and subhalo abundance matching \citep[e.g.][]{Kravtsov04}- —would produce similar mock catalogues, although an advantage of the SS09 approach is that it includes a strongly constrained model of galaxy colours."
+ S06 describes how the luminosities and real-space and redshift-space galaxy positions are modelled., S06 describes how the luminosities and real-space and redshift-space galaxy positions are modelled.
+ Our model distinguishes between the ‘central’ galaxy in a halo and all the other galaxies (‘satellites’)., Our model distinguishes between the `central' galaxy in a halo and all the other galaxies (`satellites').
+ We assume that central galaxies have the same positions and velocities as the haloes in the dark matter simulation., We assume that central galaxies have the same positions and velocities as the haloes in the dark matter simulation.
+" In other words, central galaxies are at the centre of the haloes, and the satellites are located around them."," In other words, central galaxies are at the centre of the haloes, and the satellites are located around them."
+" An important assumption in the model is that all galaxy properties—their numbers, spatial distributions, velocities, luminosities, and colours—are determined by halo mass alone."," An important assumption in the model is that all galaxy properties—their numbers, spatial distributions, velocities, luminosities, and colours—are determined by halo mass alone."
+" These galaxy properties are constrained by SDSS observations, including the luminosity function (Blantonetal.2003),, luminosity-dependent two-point clustering (Zehavietal.2005;Skibbaetal.2006;Zheng,Coil&Zehavi 2007),, and the colour-magnitude distribution and colour-dependent clustering (Skibba2009).."," These galaxy properties are constrained by SDSS observations, including the luminosity function \citep{Blanton03}, luminosity-dependent two-point clustering \citep{Zehavi05, Skibba06, Zheng07}, and the colour-magnitude distribution and colour-dependent clustering \citep{Skibba09a}."
+" Note that the clustering constraints result in a mock catalogue that approximately reproduces the observed environmental dependence of luminosity and colour, on scales of 100h!kpc to 30Mpc."," Note that the clustering constraints result in a mock catalogue that approximately reproduces the observed environmental dependence of luminosity and colour, on scales of $100\,h^{-1} {\rm kpc}$ to $30\,h^{-1} {\rm Mpc}$."
+ The number of satellite galaxies in the model follows a Poisson distribution with a mean value that increases with halo mass., The number of satellite galaxies in the model follows a Poisson distribution with a mean value that increases with halo mass.
+" The satellites are distributed around the halo centre so that they follow a Navarro,Frenk&White(1996) profile with the mass-concentration relation from Macció,Dutton&vandenBosch(2008).", The satellites are distributed around the halo centre so that they follow a \citet{Navarro96} profile with the mass-concentration relation from \citet{Maccio08}.
+. We assign redshift-space coordinates to the mock galaxies assuming that a galaxy's velocity is given by the sum of the velocity of its parent halo plus a virial motion contribution that is drawn from a Maxwell-Boltzmann distribution with dispersion that depends on halo mass (S06)., We assign redshift-space coordinates to the mock galaxies assuming that a galaxy's velocity is given by the sum of the velocity of its parent halo plus a virial motion contribution that is drawn from a Maxwell-Boltzmann distribution with dispersion that depends on halo mass (S06).
+" We specify a minimum r-band luminosity for the galaxies in the catalogue, M;—5log(h)= —19, to stay well above the resolution limit of the Millennium Simulation, avoiding any issues of completeness that may bias our results."," We specify a minimum $r$ -band luminosity for the galaxies in the catalogue, $M_r-5{\rm log}(h)=-19$, to stay well above the resolution limit of the Millennium Simulation, avoiding any issues of completeness that may bias our results."
+" We generate luminosities for the central galaxies, while accounting for the stochasticity between their luminosities and host halo mass, and then we generate the satellite luminosities so that the observed luminosity distribution is reproduced for ΛΜ,—5log(h)<—19 (S06)."," We generate luminosities for the central galaxies, while accounting for the stochasticity between their luminosities and host halo mass, and then we generate the satellite luminosities so that the observed luminosity distribution is reproduced for $M_r-5{\rm log}(h)\leq-19$ (S06)."
+" We model the observed g—r colour distribution at a given luminosity as the sum of two Gaussian components, commonly referred to as the ‘blue cloud’ and ‘red sequence’."," We model the observed $g-r$ colour distribution at a given luminosity as the sum of two Gaussian components, commonly referred to as the `blue cloud' and `red sequence'."
+" Our colour model has five constraints as a function of luminosity: the mean and scatter of the red sequence, mean and scatter of the blue cloud, and the blue fraction."," Our colour model has five constraints as a function of luminosity: the mean and scatter of the red sequence, mean and scatter of the blue cloud, and the blue fraction."
+" We assume that the colour distribution at fixed luminosity is approximately independent of halo mass, and that the satellite colour distribution varies such that its mean increases with luminosity (i.e., the satellite red fraction increases with luminosity in a particular way)."," We assume that the colour distribution at fixed luminosity is approximately independent of halo mass, and that the satellite colour distribution varies such that its mean increases with luminosity (i.e., the satellite red fraction increases with luminosity in a particular way)."
+ These two assumptions are tested and verified with galaxy group catalogues in Skibba(2009).., These two assumptions are tested and verified with galaxy group catalogues in \citet{Skibba09a}.
+" This procedure produces a mock galaxy catalogue containing 1.84 million galaxies, of which 29 percent are satellites."," This procedure produces a mock galaxy catalogue containing 1.84 million galaxies, of which 29 percent are satellites."
+ Galaxies occupy haloes with masses ranging from 1011 to 11015:3ΗMo., Galaxies occupy haloes with masses ranging from $10^{11}$ to $10^{15.3}h^{-1} M_\odot$.
+ We also construct a mock light cone from the catalogue by selecting galaxies that are within a radial distance of 500!Mpc from one corner of the box.," We also construct a mock light cone from the catalogue by selecting galaxies that are within a radial distance of $500\,h^{-1} {\rm Mpc}$ from one corner of the box."
+" This gives an opening angle of 90x degrees and a depth of 500A!Mpc, for which right ascension and declinations are determined."," This gives an opening angle of $90\times90$ degrees and a depth of $500\,h^{-1} {\rm Mpc}$, for which right ascension and declinations are determined."
+The analysisinSection 4 is carried out using a sample of galaxies that are common to both the boxandthe cone and are chosen so not to be affected by edges.,The analysisinSection \ref{sec:res} is carried out using a sample of galaxies that are common to both the boxandthe cone and are chosen so not to be affected by edges.
+ Figure shows the mean number of galaxies as a function of halo, Figure \ref{fig:pop} shows the mean number of galaxies as a function of halo
+and influences.,and influences.
+ 118325-5926 was observed on 2002 March 20 (Obs ID 3148) and on 2002 March 23 (Obs ID 3452) with the ACIS-S (Advanced CCD Imaging Spectrometer - Spectral component) instrument coupled with the HETGS onChandra., 18325-5926 was observed on $2002$ March $20$ (Obs ID $3148$ ) and on $2002$ March $23$ (Obs ID $3452$ ) with the ACIS-S (Advanced CCD Imaging Spectrometer - Spectral component) instrument coupled with the HETGS on.
+".. The exposure times of the two observations were 56.9 ks and 51.09 ks respectively,", The exposure times of the two observations were $56.9$ ks and $51.09$ ks respectively.
+ The light-curve count rate of the dispersed. data fluctuated between a low of ~0.1 cts ! to a high of ~0.8 cis ! (Fig. 1))., The light-curve count rate of the dispersed data fluctuated between a low of $\sim0.1$ cts $^{-1}$ to a high of $\sim0.8$ cts $^{-1}$ (Fig. \ref{cntsT}) ).
+ The data were reprocessed with CIAO v 4.1.2 and new level 2 event [iles were created., The data were reprocessed with CIAO v $4.1.2$ and new level $2$ event files were created.
+ The +1 and —] orders of the WEG and MEG erating arms were combined. as were both the March. 20 and 23 observations.," The $+1$ and $-1$ orders of the HEG and MEG grating arms were combined, as were both the March $20$ and $23$ observations."
+ For Fe emission line analvsis. we also investigated the observations taken a vear earlier.," For Fe emission line analysis, we also investigated the observations taken a year earlier."
+ Data was binned by a factor of 4 (IETGCS resolution)., Data was binned by a factor of $4$ (HETGS resolution).
+ There are enough counts to decisively study. absorption lines in the 1.0 7.0 keV spectral band., There are enough counts to decisively study absorption lines in the $1.0$ $7.0$ keV spectral band.
+ This region contains counts between 5 and 40 per bin., This region contains counts between $5$ and $40$ per bin.
+ The signal-to-noise ratio is al least 3:1 in this region. bul is much lower outside (his range.," The signal-to-noise ratio is at least $3:1$ in this region, but is much lower outside this range."
+ We analvze the MEG spectrum from 0.8. 1.2 keV. and the VEG spectrum from 1.2 8.0 keV (extending sliehtlv into the reeions of lower counts). and we cross-check absorption features in the IEG and MEG spectra between 1.2 2.0 keV. Low counts («5 counts per bin) are seen in the MEG data from 0.8 1.0 keV and in the WEG data from 1.2 1.3 keV and 7.0 8.0 keV. We use the Interactive Spectral Interpretation System. (ISIS) {η} (o fit the spectra.," We analyze the MEG spectrum from $0.8$ $1.2$ keV, and the HEG spectrum from $1.2$ $8.0$ keV (extending slightly into the regions of lower counts), and we cross-check absorption features in the HEG and MEG spectra between $1.2$ $2.0$ keV. Low counts $<5$ counts per bin) are seen in the MEG data from $0.8$ $1.0$ keV and in the HEG data from $1.2$ $1.3$ keV and $7.0$ $8.0$ keV. We use the Interactive Spectral Interpretation System (ISIS) \citep{2002hrxs.confE..17H} to fit the spectra."
+ Absorption features are found by fitting the spectrum with an ionized (warm) absorber. WARMABS. using the photoionization code (?)..," Absorption features are found by fitting the spectrum with an ionized (warm) absorber, , using the photoionization code \citep{2001ApJS..134..139B}."
+ The ionization parameter £. column density. and outflow: velocity are variable parameters in (he model which are optimized [or best-fit.," The ionization parameter $\xi$, column density, and outflow velocity are variable parameters in the model which are optimized for best-fit."
+ All spectral fits account for the line-ol-sight Galactic column Ay=7.4x107 em7., All spectral fits account for the line-of-sight Galactic column $N_{\rm H}=7.4\times 10^{20} $ $^{-2}$.
+ The model (72) which corrects for X-ray absorption due to gas-phase and grain-phase ISM. and molecules in the ISM. is applied for this purpose.," The model \citep{2006HEAD....9.1360W} which corrects for X-ray absorption due to gas-phase and grain-phase ISM, and molecules in the ISM, is applied for this purpose."
+ A gap in the detector chip. where the effective area is small. was noticed in the +1 order of the HEG erating arm in both observations of 118325-5926.," A gap in the detector chip, where the effective area is small, was noticed in the $+1$ order of the HEG grating arm in both observations of 18325-5926."
+ The —1 grating arn appeared normal., The $-1$ grating arm appeared normal.
+ The spectral baud affected is between 2.5 2.7 keV. This region is nol reliable for the detection of absorption leatures in our source. and (therefore is ignored [or," The spectral band affected is between $2.5$ $2.7$ keV. This region is not reliable for the detection of absorption features in our source, and therefore is ignored for"
+"[rom 11.2 eV to 13.6 eV. which vields ~107"" erg s+ τσ tse +.","from 11.2 eV to 13.6 eV, which yields $\sim 10^{-20}$ erg $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$."
+" Thus. for a general radiation field the LED photocissociation time can be expressed. as fais. ~ 107 ve CJu £10.4) 1; where Jp measures the LAV flux in units of 7 erg s bem? bse, "," Thus, for a general radiation field the HD photodissociation time can be expressed as $t_{\rm diss}$ $\sim$ $^8$ yr $J_{\rmn LW}$ $^{-4}$ $^{-1}$, where $J_{\rmn LW}$ measures the LW flux in units of $^{-21}$ erg $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$."
+"To estimate the importance of radiation emitted. by shocked primordial gas in destroving LD molecules. we compare the LID photodissociation timescale. fai... to the formation timescale. fia, (see Machida et al."," To estimate the importance of radiation emitted by shocked primordial gas in destroying HD molecules, we compare the HD photodissociation timescale, $t_{\rm diss}$, to the formation timescale, $t_{\rm form}$ (see Machida et al."
+ 2005)., 2005).
+ Phe LW photon lux emergent from a 100 kms + shock passing through a medium. of pre-shock hydrogen number density of 10 is ~107 photons em7 I. which corresponds to Jp~10I (Shull Silk 1979).," The LW photon flux emergent from a 100 km $^{-1}$ shock passing through a medium of pre-shock hydrogen number density of $10^2$ $^{-3}$ is $\sim
+10^2$ photons $^{-2}$ $^{-1}$, which corresponds to $J_{LW} \sim 10^{-4}$ (Shull Silk 1979)."
+ This results in [ai ~ QU ve (Le Petit. Ποιο Le Bourlot 2002).," This results in $t_{\rm diss}$ $\sim$ $^8$ yr (Le Petit, Roueff Le Bourlot 2002)."
+" From our calculations. we estimate /p,44,10"" vr for the sane shock velocity. and pre-shock density."," From our calculations, we estimate $t_{\rm form} \sim 10^{6}$ yr for the same shock velocity and pre-shock density."
+ ας. for cases in which the post-shock gas is the only source of LAW photons. he formation timescale is shorter than the. dissociation imeseale by two orders of magnitude. and the clleet of ohotocdissociating radiation produced by the shock is likely. o be unimportant over the range of shock velocities and e-shock densities that we consider in this paper.," Thus, for cases in which the post-shock gas is the only source of LW photons, the formation timescale is shorter than the dissociation timescale by two orders of magnitude, and the effect of photodissociating radiation produced by the shock is likely to be unimportant over the range of shock velocities and pre-shock densities that we consider in this paper."
+ We note. however. that in relic LL HE. regions the ohotodestruction. of LID may be particularly important due to emission of LAV radiation through the two-photon decav of recombining helium atoms. at densities. below ~ 10t  (Mathis. 1957: Pottasch 1901).," We note, however, that in relic H II regions the photodestruction of HD may be particularly important due to emission of LW radiation through the two-photon decay of recombining helium atoms, at densities below $\sim$ $^4$ $^{-3}$ (Mathis 1957; Pottasch 1961)."
+ While helium recombination radiation will also be generated by shocks strong enough to ionize helium. the column densities of ionized helium in the first relic LE 1E regions are expected o be much higher than those in shock fronts. since the irst Hl HE regions generated by massive Pop LLL stars are expected to have been of order a few kpe in diameter with the majority of the helium being tonizecl by the iud UV emission from the central star (Ixitavana et al.," While helium recombination radiation will also be generated by shocks strong enough to ionize helium, the column densities of ionized helium in the first relic H II regions are expected to be much higher than those in shock fronts, since the first H II regions generated by massive Pop III stars are expected to have been of order a few kpc in diameter with the majority of the helium being ionized by the hard UV emission from the central star (Kitayama et al."
+ 2004: Whalen. Abel Norman 2004: Alvarez. Dromm Shapiro 2005).," 2004; Whalen, Abel Norman 2004; Alvarez, Bromm Shapiro 2005)."
+ With these larger column. densities. the luxes of recombination radiation from helium in relic LE LE regions may be strong enough to ellicientlv. photodissociate LD imoleleules.," With these larger column densities, the fluxes of recombination radiation from helium in relic H II regions may be strong enough to efficiently photodissociate HD molelcules."
+ However. if molecule formation takes place elliciently after the gas has cooled below the temperatures at which helium recombines. then the elect of the destruction of molecules by helium: recombination radiation could. be mitigated," However, if molecule formation takes place efficiently after the gas has cooled below the temperatures at which helium recombines, then the effect of the destruction of molecules by helium recombination radiation could be mitigated."
+ A careful treatment of the ellect of recombination racliation is clearly required in order to accurately caleulate 1ο chemical and thermal evolution of relie LE LE regions., A careful treatment of the effect of recombination radiation is clearly required in order to accurately calculate the chemical and thermal evolution of relic H II regions.
+ We defer such a treatment to future work., We defer such a treatment to future work.
+ If the column density of LID is high enough that πο-shielding is important. the value of fai. caleulated here will represent a conservative lower limit on the photoclissociation time of shielded LID molecules.," If the column density of HD is high enough that self-shielding is important, the value of $t_{\rm diss}$ calculated here will represent a conservative lower limit on the photodissociation time of shielded HD molecules."
+ The ellects. of Ils selt-shielding have been widely studied and have been found to have important consequences for the cooling of primordial gas subject to a LW background radiation field (Draine Dertoldi 1996: HLaiman et al., The effects of $_2$ self-shielding have been widely studied and have been found to have important consequences for the cooling of primordial gas subject to a LW background radiation field (Draine Bertoldi 1996; Haiman et al.
+ 1997: Yoshida et al., 1997; Yoshida et al.
+ 2003)., 2003).
+ As well. it has been shown that the dissociation rate of LID shielded by Ls is expected to be lower than the dissociation rate of unshieldded LID by up to a factor of about 1/3. due to the close overlapping of several LID lines by H2 lines (Barsuhn 1977).," As well, it has been shown that the dissociation rate of HD shielded by $_2$ is expected to be lower than the dissociation rate of unshieldded HD by up to a factor of about 1/3, due to the close overlapping of several HD lines by $_2$ lines (Barsuhn 1977)."
+ Here. for simplicity. we neglect the effects of both LID selt-shielding and LED shielding by 1H».," Here, for simplicity, we neglect the effects of both HD self-shielding and HD shielding by $_2$."
+ We note. iowever. that if there exists a &eneral LAW background in addition to the radiation emitted locally by the shocked σας. such as from first generation massive stars during structure ormation or from continuum radiation from recombining wlium atoms in relic Hb LL regions. then sel(-shielding ellects may become much more important in preventing the ohotocissociation of LIL).," We note, however, that if there exists a general LW background in addition to the radiation emitted locally by the shocked gas, such as from first generation massive stars during structure formation or from continuum radiation from recombining helium atoms in relic H II regions, then self-shielding effects may become much more important in preventing the photodissociation of HD."
+ The cooling properties of LID are distinct from those of Ll» for the Following reasons (e.g. Flower et al., The cooling properties of HD are distinct from those of $_2$ for the following reasons (e.g. Flower et al.
+ 2000: Uchara Inutsuka 2000)., 2000; Uehara Inutsuka 2000).
+ Firstly. LID has a permanent dipole moment. allowing dipole rotational transitions. which spontaneously occur much more often. than the quadrupole: rotational transitions in Hl».," Firstly, HD has a permanent dipole moment, allowing dipole rotational transitions, which spontaneously occur much more often than the quadrupole rotational transitions in $_2$."
+ Also. the dipole moment of HD. allows rotational transitions of AJ=ΕΙ. which are of lower energy than the AZ=+2 quadrupole transitions of Ls.," Also, the dipole moment of HD allows rotational transitions of $\Delta J = \pm 1$, which are of lower energy than the $\Delta J = \pm 2$ quadrupole transitions of $_2$."
+ Thus. collisional excitation of LID from the erouncd to the first excited rotational level (/=1). and the subsequent radiative decay back to the ground state by a dipole transition. can allow HD to elliciently cool gas to lower temperatures than can be reached with Ll» line cooling alone.," Thus, collisional excitation of HD from the ground to the first excited rotational level $J$ =1), and the subsequent radiative decay back to the ground state by a dipole transition, can allow HD to efficiently cool gas to lower temperatures than can be reached with $_2$ line cooling alone."
+ Finally. the rotational energy levels of HD. are lower than those of LH» by a factor of the reduced mass. which is higher by a factor of ~4/3 lor LID than for LH». allowing for even lower energy collisions to excite the first rotational level.," Finally, the rotational energy levels of HD are lower than those of $_2$ by a factor of the reduced mass, which is higher by a factor of $\sim 4/3$ for HD than for $_2$, allowing for even lower energy collisions to excite the first rotational level."
+ We have implemented the LED cooling function provided by Flower et al. (, We have implemented the HD cooling function provided by Flower et al. (
+2000: see also Lipovka ct al.,2000; see also Lipovka et al.
+ 2005)., 2005).
+ To take into account that the gas cannot radiatively cool to below the temperature of the cosmic microwave background (CM). we employ an ellective cooling rate. according to where Zo... is the temperature of the gas.," To take into account that the gas cannot radiatively cool to below the temperature of the cosmic microwave background (CMB), we employ an effective cooling rate, according to where $T_{\rmn gas}$ is the temperature of the gas."
+ For a gas which cools purely radiatively. and πο by acliabatic expansion. the CAIB sets a lower limit on the temperature to which it can cool (c.g. Larson 1998).," For a gas which cools purely radiatively, and not by adiabatic expansion, the CMB sets a lower limit on the temperature to which it can cool (e.g. Larson 1998)."
+ To show this. we will consider the process by which the primordial gas cools to temperatures approaching the CMD value.," To show this, we will consider the process by which the primordial gas cools to temperatures approaching the CMB value."
+ The frequeney of emitted. radiation for the rotational transition J=1 ο) is given by where £o is the frequency for the J=1.+0 rotational transition. Aj is the Boltzmann constant. and P is the Planck," The frequency of emitted radiation for the rotational transition $J = 1 \to 0$ of HD is given by where $\nu_{10}$ is the frequency for the $J=1 \to 0$ rotational transition, $k_{\rmn B}$ is the Boltzmann constant, and $h$ is the Planck"
+the tvpe of reaction that takes place when a self-consistent. though simple. recipe for the injection of particles from. the thermal pool is adopted.,"the type of reaction that takes place when a self-consistent, though simple, recipe for the injection of particles from the thermal pool is adopted."
+ Εις recipe is similar to that proposed bv Kane. Jones Cieseler. (2002) in terms of the underlying physical interpretation of the injection. but probably simpler in its implementation.," This recipe is similar to that proposed by Kang, Jones Gieseler (2002) in terms of the underlying physical interpretation of the injection, but probably simpler in its implementation."
+" For non-relativistic shocks. the cistribution of particles downstream is quasi-isotropic. so that the [lux of particles crossing the shock surface from downstream to upstream can be written as where e(p) is the velocity of particles with momentum p and mis the shock speed in the frame comoving with the ownstream Εις,"," For non-relativistic shocks, the distribution of particles downstream is quasi-isotropic, so that the flux of particles crossing the shock surface from downstream to upstream can be written as where $v(p)$ is the velocity of particles with momentum $p$ and $u_d$ is the shock speed in the frame comoving with the downstream fluid."
+ Phe term m;|ep) is the component dong the direction. perpendicular to the shock surface. of 10 velocity of particles with momentum p moving in the irection fr., The term $u_d+v(p)\mu$ is the component along the direction perpendicular to the shock surface of the velocity of particles with momentum $p$ moving in the direction $\mu$.
+ Lt follows that the flux of particles moving angent to the shock surface (namely with (i—mafelp)) is zero., It follows that the flux of particles moving tangent to the shock surface (namely with $\mu=-u_d/v(p)$ ) is zero.
+ We recall that. having in mind. collisionless shocks. 10 twpical thickness of the shock. A. is the collision length ssociated with the magnetic interactions that give rise to 1ο formation of the discontinuity.," We recall that, having in mind collisionless shocks, the typical thickness of the shock, $\lambda$, is the collision length associated with the magnetic interactions that give rise to the formation of the discontinuity."
+ Useless to sav that these interactions are all but well. known. and at present. the yest we can do is to attempt a phenomenological approach o take them into account. without having to deal with heir detailed: physical. understanding.," Useless to say that these interactions are all but well known, and at present the best we can do is to attempt a phenomenological approach to take them into account, without having to deal with their detailed physical understanding."
+ His however worth recalling that many attempts have been mace to tackle the oblem of injection at a more fundamental level (Malkoy 1998: Malkov Vollk 1995: Alalkov Wollk 1998)., It is however worth recalling that many attempts have been made to tackle the problem of injection at a more fundamental level (Malkov 1998; Malkov Völlk 1995; Malkov Völlk 1998).
+" Lore. we consider. the reasonable. situation.⋠⋠ in which. ÀXry’.Hu where rj""Hxpis ds. the Larmor radius. of⋅ the particles. in. he downstream Iluid that carry most of the thermal energy. namely those with: momentum 1.5pi, (pi,=(δημο).qual2 rere is the momentum of the particles in the thermal peas of the maxwellian clistribution in the downstream. plasma. jwing temperature 75)."," Here, we consider the reasonable situation in which $\lambda\propto r_L^{th}$, where $r_L^{th}\propto p_{th}$ is the Larmor radius of the particles in the downstream fluid that carry most of the thermal energy, namely those with momentum $1.5~p_{th}$ $p_{th}=(2 m k_BT_2)^{1/2}$ here is the momentum of the particles in the thermal peak of the maxwellian distribution in the downstream plasma, having temperature $T_2$ )."
+ We stress here the important point hat the temperature of the downstream gas (and therefore pin) ds determined. by the shock strength. which in the oesence of accelerated. particles. is alfected by the pressure of the non-thermal component.," We stress here the important point that the temperature of the downstream gas (and therefore $p_{th}$ ) is determined by the shock strength, which in the presence of accelerated particles, is affected by the pressure of the non-thermal component."
+ In particular. the higher the ellicieney of the shock as à particle accelerator. the weaker --s ellicienev in terms of heating of the background plasma (sec section NJ) ," In particular, the higher the efficiency of the shock as a particle accelerator, the weaker its efficiency in terms of heating of the background plasma (see section \ref{sec:heating}) )."
+For collisionless shocks. it is not clear whether the ownstream plasma can actually be thermalized. and. the istribution function be a maxwellian.," For collisionless shocks, it is not clear whether the downstream plasma can actually be thermalized and the distribution function be a maxwellian."
+ On the other hand. -- is generally assumed. that this is the case. so that. in 1e following we consider the case in which the bulk of 16 background. plasma is thermal ancl has a maxwellian μα»eetrum. at. temperature Z given by the generalized tankine-Llugoniot relations in the presence of accelerated xwiicles (see section 2)).," On the other hand, it is generally assumed that this is the case, so that in the following we consider the case in which the bulk of the background plasma is thermal and has a maxwellian spectrum at temperature $T$ given by the generalized Rankine-Hugoniot relations in the presence of accelerated particles (see section \ref{sec:nonlin}) )."
+ For modified shocks. the points iscussecl above apply to the so-called subshock. where the injection of particles from the thermal pool is expected to ake place.," For modified shocks, the points discussed above apply to the so-called subshock, where the injection of particles from the thermal pool is expected to take place."
+ We recall that for strongly. modified. shocks the subshock is weak. and rather ineLllicient in the heating of the xckeround. plasma.," We recall that for strongly modified shocks the subshock is weak, and rather inefficient in the heating of the background plasma."
+ From Eq., From Eq.
+ 22. we get:, \ref{eq:flux1} we get:
+The cdillerences. between the late reionization. moclels in comparison with the expected: sensitivity. of the mission can be expressed in terms of the power spectrum Ον) (for. the anisotropy. £ and TE component of polarization) where the indices / and j denote the different. models.,"The differences between the late reionization models in comparison with the expected sensitivity of the mission can be expressed in terms of the power spectrum $C_p (\ell)$ (for the anisotropy, $E$ and $TE$ component of polarization) where the indices $i$ and $j$ denote the different models."
+ In order to clarify the manifestations of the complex ionization regimes in the models 1 and 2 we need to compare the peak to peak amplitudes of the D;(0) function with the expected. error of the anisotropy. power spectrum for the experiment.," In order to clarify the manifestations of the complex ionization regimes in the models 1 and 2 we need to compare the peak to peak amplitudes of the $D_{i,j}(\ell)$ function with the expected error of the anisotropy power spectrum for the experiment."
+ We assume that the systematics and Foreground effects are successtully removed., We assume that the systematics and foreground effects are successfully removed.
+" The corresponding error bar should be where w=(0,,6pywHar) Weceexp[(0|0/26]. fas&0.65 is the sky coverage during the first vear of observations. o, is the sensitivity per resolution element ÜrwHMÜpw.par add; (£u—=y/Sin?2g.6us F"," The corresponding error bar should be where $w=(\sigma_{p}\theta_{\rm FWHM})^{-2}$, $W_{\ell}\simeq
+\exp\left[-\ell(\ell+1)/2\ell^2_s\right]$, $f_{\rm sky} \simeq 0.65$ is the sky coverage during the first year of observations, $\sigma_{p}$ is the sensitivity per resolution element $\theta_{\rm FWHM} \times \theta_{\rm FWHM}$ and $\ell_s=\sqrt{8\ln 2}
+\, \theta^{-1}_{\rm FWHM}$."
+or all frequeney channels. for example. the JFNLIM. are less than 30 aremin. so for the estimation of he errors at {40 range we can omit the second term in Eq. (10)).," For all frequency channels, for example, the FWHM are less than 30 arcmin, so for the estimation of the errors at $\ell
+\le 40$ range we can omit the second term in Eq. \ref{c}) )."
+ In Fig.3 we show the polarization power spectrum and “PE cross-correlation for modelsο, In Fig.3 we show the polarization power spectrum and TE cross-correlation for models.
+ς While polarization power spectrum is not observed by WALAPR. the upcoming tanek mission will be able to provide us the observation o differentiate different models.," While polarization power spectrum is not observed by WMAP, the upcoming Planck mission will be able to provide us the observation to differentiate different models."
+ Lt is clear from Fig that he PE cross-correlation spectra for dillerent regimes of reionization have dillerent. shape at f10 range but do not vary significantly for £2710., It is clear from Fig.4 that the TE cross-correlation spectra for different regimes of reionization have different shape at $\ell \le 10$ range but do not vary significantly for $ \ell > 10$.
+ ALL the deviations Lie inside he cosmic variance and are practically not observable., All the deviations lie inside the cosmic variance and are practically not observable.
+ As one can see from Fig., As one can see from Fig.
+ 5. for {διο} the corresponding peak to peaks amplitudes are on the order of magnitude 20% ," \ref{diff} for $
+D_{1,2}(\ell)$ the corresponding peak to peak amplitudes are on the order of magnitude $20\%$ "
+covariance provide independence on OT. ancl wind: speed vertical profile.,covariance provide independence on OT and wind speed vertical profile.
+ From the expression. (33)) it follows that £07). differs from square averaged. wind 1»5 in additional multiplicand 217(0)/C which varies by 23 times depending on wind direction in the case of /- and /-motions., From the expression \ref{eq:v2_def}) ) it follows that $\langle v^2\rangle$ differs from square averaged wind $\bar V_2$ in additional multiplicand $2{\mathcal L}^*(\theta)/{\mathcal L}$ which varies by 2–3 times depending on wind direction in the case of $t$ - and $l$ -motions.
+ At the same time. for the c-motion this factor deviates from 1 by no more than 5 per cent (see Table 1)) and angular. dependence can be neglected.," At the same time, for the $c$ -motion this factor deviates from 1 by no more than 3 per cent (see Table \ref{tab:Lint}) ) and angular dependence can be neglected."
+ Therefore. in this case. one can assume NEN7 exactly.," Therefore, in this case, one can assume $\bar V_2 = \langle v^2\rangle^{1/2}$ exactly."
+ The definition of 9e and the expression (37)) set the connection between normalized correction for. covariance and measured correlation coellicient. jx what allows to express square averaged. wind speed. Y» through instrumental parameters and measured p: This clepencdence gives. reasonably σου estimation of the 45 while quadratic approximation for de(a) is valid., The definition of $\delta c$ and the expression \ref{eq:corr_fin}) ) set the connection between normalized correction for covariance and measured correlation coefficient $\tilde\rho$: what allows to express square averaged wind speed $\bar V_2$ through instrumental parameters and measured $\tilde\rho$: This dependence gives reasonably good estimation of the $\bar V_2$ while quadratic approximation for $\delta c(\hat\omega)$ is valid.
+ For greater wind speeds. rise of def) slows down and evaluated Ve becomes underestimated.," For greater wind speeds, rise of $\delta c(\hat\omega)$ slows down and evaluated $\bar V_2$ becomes underestimated."
+ This οσοι is illustrated in Fig. 5.., This effect is illustrated in Fig. \ref{fig:wind}.
+ For our DIAIAIL device substitution of numerical values in formula (40)) gives 1»=13.5(1p27.," For our DIMM device substitution of numerical values in formula \ref{eq:wind_res}) ) gives $\bar V_2=13.8\,(1-\tilde\rho)^{1/2}$."
+ Corresponding dependence is shown in figure by dashed Line., Corresponding dependence is shown in figure by dashed line.
+ Comparison of the exact dependence with approximate shows that short. exposure. approximation underestimates wind speed by 13 por cent and 32 per cont forSms+ and lOnis+. respectively.," Comparison of the exact dependence with approximate shows that short exposure approximation underestimates wind speed by 13 per cent and 32 per cent for$5\rmn{\,m\,s^{-1}}$ and $10\rmn{\,m\,s^{-1}}$, respectively."
+ Exposure time has almost. no elLect on appearance of this curve: on the contrary. change in the period 1) leads to proportional change of y axis scale.," Exposure time has almost no effect on appearance of this curve; on the contrary, change in the period $T$ leads to proportional change of $y$ axis scale."
+ eaturallv. impact of fast turbulence lavers will be significantIy weakened.," Naturally, impact of fast turbulence layers will be significantly weakened."
+ Therefore. the wind speed. derived with DIMAM can be considered only as lower estimate of the 1».," Therefore, the wind speed derived with DIMM can be considered only as lower estimate of the $\bar V_2$."
+ To do it more realistic. CCD camera frame period 1 should. be reduced twice at least.," To do it more realistic, CCD camera frame period $T$ should be reduced twice at least."
+ Phen. the underestimate zz30 per cent will be provided for turbulent. laver. with 20ms* wind.," Then, the underestimate $\approx 30$ per cent will be provided for turbulent layer with $20\rmn{\,m\,s^{-1}}$ wind."
+ Additionally. certain improvement. can be achieved by increase of DIMM aperture. 2.," Additionally, certain improvement can be achieved by increase of DIMM aperture $D$."
+ Nevertheless. the method is able to give quite good estimates even with existent device. as will be shown in Sect.," Nevertheless, the method is able to give quite good estimates even with existent device, as will be shown in Sect."
+ 6.3. on real data., \ref{sec:MASSwind} on real data.
+ Considered average wind speed Vo dilfers slightly. [rom το entering the definition of atmospheric coherence time Ty., Considered average wind speed $\bar V_2$ differs slightly from $\bar V_{5/3}$ entering the definition of atmospheric coherence time $\tau_0$.
+ Nevertheless. investigation bv ? has shown that 15 value may be used for estimation of το.," Nevertheless, investigation by \citet{TKel2007} has shown that $\bar V_2$ value may be used for estimation of $\tau_0$."
+ Recall that method of derivation of 75; implemented in EXDIZ (Fast. Defocusing of stellar image) (72). and MASS instrument (2). are based on the V5 either., Recall that method of derivation of $\tau_0$ implemented in FADE (Fast Defocusing of stellar image) \citep{FADE2008} and MASS instrument \citep{tau2011} are based on the $\bar V_2$ either.
+ Experimental proof of the basic relations was mace using ALASS/DIMM data obtained on Mount. Shatdjatmaz in 20072010 (?) for characterisation of OT., Experimental proof of the basic relations was made using MASS/DIMM data obtained on Mount Shatdjatmaz in 2007–2010 \citep{kgo2010} for characterisation of OT.
+ The program records measurements results in output file every 2 s and also its l-münute averages., The program records measurements results in output file every 2 s and also its 1-minute averages.
+ These results include variances ear ancl covariances of the dilferential image motions and coordinate noise ear? in ongitucdinal and transverse directions., These results include variances $var$ and covariances of the differential image motions and coordinate noise $var^{*}$ in longitudinal and transverse directions.
+ Value of coordinate noise is computed during processing of frames using formula (S) from (7) which gives slightly. underestimated: values or thresholding centroiding method since it does not take into account [Iuctuations of number of pixels exceeding a hresholel., Value of coordinate noise is computed during processing of frames using formula (8) from \citep{MD2007} which gives slightly underestimated values for thresholding centroiding method since it does not take into account fluctuations of number of pixels exceeding a threshold.
+ Though in the real conditions. impact of the noise is very small (median of ratio car’fearzz0.0008 and in 99.7 »r cent of cases it is less than 0.01). nevertheless. it is aken into account in further processing in obvious way: 6?—varrar’.," Though in the real conditions, impact of the noise is very small (median of ratio $var^{*}/var \approx 0.0008$ and in 99.7 per cent of cases it is less than $0.01$ ), nevertheless, it is taken into account in further processing in obvious way: $\tilde\sigma^2 = var - var^{*}$."
+ Covariance is not biased by the noise but incorrect account for the noise in 67 can lead to bias in the correlation cocllicient., Covariance is not biased by the noise but incorrect account for the noise in $\tilde\sigma^2$ can lead to bias in the correlation coefficient.
+ Distributions of the correlation coellicients in 20072009 (containing z90000 l-minute points) have similar characteristics for all image motions: medians are equal to 0.85. only for 4 per cent of events p«0.6 and for 1 per cent p< OAL.," Distributions of the correlation coefficients in 2007--2009 (containing $\approx 90\,000$ 1-minute points) have similar characteristics for all image motions: medians are equal to $0.85$, only for 4 per cent of events $\tilde\rho < 0.6$ and for 1 per cent — $\tilde\rho < 0.44$ ."
+ On the other hand. only for 5 per cent of the measurements. ο>0.95 and 0.5 per cent 2 0.97.," On the other hand, only for 5 per cent of the measurements, $\tilde\rho > 0.95$ and 0.5 per cent $>0.97$ ."
+ Distributions of 2-s points demonstrate similar. behavior slightly different in details., Distributions of 2-s points demonstrate similar behavior slightly different in details.
+ Prior to theoretical analysis described. in this paper. during processing NLASS/DIMM data in 2010. we used the following empirical law (basedon numerical calculations. see Sect. 3.1):," Prior to theoretical analysis described in this paper, during processing MASS/DIMM data in 2010, we used the following empirical law (basedon numerical calculations, see Sect. \ref{sec:sp_filters}) ):"
+Tu |?] three different methods were developed to solve the boundary value problem(56).,In \cite{Cook-Choptuik-etal:1993} three different methods were developed to solve the boundary value problem.
+. The best one turned out to be a finite difference code based ou a specially adapted coordinate svsteim. the so-called Cacdézz coordinates.," The best one turned out to be a finite difference code based on a specially adapted coordinate system, the so-called Cadézz coordinates."
+ This code is an FAS-inultierid aleorithui developed specifically for the differenti: equation(53)., This code is an FAS-multigrid algorithm developed specifically for the differential equation.
+. Care was taken that the truncation error is strictly even in oOovid-spacing h. thus allowing one to take two or three solutions at different resolutious aud Bichirdsou extrapolate to.»0.," Care was taken that the truncation error is strictly even in grid-spacing $h$, thus allowing one to take two or three solutions at different resolutions and Richardson extrapolate to $h\to 0$."
+ The Cadézz code is thus specially built for this equation in this ecometry aud it is uulikelv that it can be sguificautlv iuproved upon by auv finite difference method., The Cadézz code is thus specially built for this equation in this geometry and it is unlikely that it can be significantly improved upon by any finite difference method.
+ Ou the other haud. our spectral solver is &eueral purpose.," On the other hand, our spectral solver is general purpose."
+ The domain decomposition is uot restricted to R with two excised spheres and we do not eimuplov auy specific optimizations for this particular problem., The domain decomposition is not restricted to $\mathbbmss{R}^3$ with two excised spheres and we do not employ any specific optimizations for this particular problem.
+ We compare these two codes for the configuration with equal sized spleres., We compare these two codes for the configuration with equal sized spheres.
+ Figure 7. shows a plot of runtime vs. achieved accuracy for both codes., Figure \ref{fig:CompareCadez} shows a plot of runtime vs. achieved accuracy for both codes.
+ These ruus were performed on a single RS6000 processor: the highest resolution runs needed about LOB of memory., These runs were performed on a single RS6000 processor; the highest resolution runs needed about 1GB of memory.
+ The solid line labeled FD represents the results of the finite difference code without Richardsou extrapolation., The solid line labeled FD represents the results of the finite difference code without Richardson extrapolation.
+ This line couverges quadratically iu erid spacing., This line converges quadratically in grid spacing.
+ The two stars represent Richardsou extrapolated, The two stars represent Richardson extrapolated
+clouds formed at airmass 2<z«10.,clouds formed at airmass $ 2 < z < 10$.
+ The equivalent. grey level plot of Fig., The equivalent grey level plot of Fig.
+ 14. for airmass less than 2 (not. shown) does not include the contribution of clouds formed. in the carly morning and late afternoon hours. specially curing the summer months.," \ref{month_hour} for airmass less than 2 (not shown) does not include the contribution of clouds formed in the early morning and late afternoon hours, specially during the summer months."
+ Using the solar radiation data recorded by the ‘PALIT site-testing group at SPAL observatory. we applied a nove method developed by Carrascoetal.(2009). to estimate the time when the sky. is clear of clouds.," Using the solar radiation data recorded by the TMT site-testing group at SPM observatory, we applied a novel method developed by \citet{Carrasco09} to estimate the time when the sky is clear of clouds."
+ From the globa normalized observed. distribution of v: the solar Εν 0). divided by the nominal solar Hux at. the. top of the atmosphere f°.cosος(E) we obtained that S2.4 per cent of the time the skv is clear of clouds for airmasses z52.," From the global normalized observed distribution of $\psi$: the solar flux $F(t)$, divided by the nominal solar flux at the top of the atmosphere $F_{\sun}\cos\theta_{\sun}(t)$ we obtained that 82.4 per cent of the time the sky is clear of clouds for airmasses $ z \le 2$."
+ This result is consistent within 6 per cent with the indepencen estimation of cloud cover using satellite data by Erasmus&VanStaden(2002) who calculated a nighttime usefu raction of SI per cent ancl estimated that the daily. clear raction is about 5 per cent less.," This result is consistent within 6 per cent with the independent estimation of cloud cover using satellite data by \citet{Erasmus02}
+ who calculated a nighttime useful fraction of 81 per cent and estimated that the daily clear fraction is about 5 per cent less."
+ By analvzing the fits to the histograms of c per month. we found that there is an annual seasonal ellect: the centre of he clear peak has a maximum of 0.889 in February. reaches a minimum of 0.761 in July and increases to 0.885 towards he end of the vear.," By analyzing the fits to the histograms of $\psi$ per month, we found that there is an annual seasonal effect: the centre of the clear peak has a maximum of 0.889 in February, reaches a minimum of 0.761 in July and increases to 0.885 towards the end of the year."
+ As an example we presented. the fits o the histograms for July ancl November obtained: using he complete data set., As an example we presented the fits to the histograms for July and November obtained using the complete data set.
+ The centre of the peak for July is ).761 and for November is 0.841 but the amount of racliation corresponding to 0=0.761 received at the site in July is 277 Wm more than that in November., The centre of the peak for July is 0.761 and for November is 0.841 but the amount of radiation corresponding to $\psi=0.761$ received at the site in July is 277 $^{-2}$ more than that in November.
+ This result. is consistent with the seasonal variation in DWV. measured at the site for several vears obtained by Otárolaοἱal.(2009) and Otárolaetal.(2010)., This result is consistent with the seasonal variation in PWV measured at the site for several years obtained by \citet{Otarola09} and \citet{Otarola10}.
+. Phe dillerence can be explained by the presence of more aerosols in July that absorb the incident raciation in the atmosphere., The difference can be explained by the presence of more aerosols in July that absorb the incident radiation in the atmosphere.
+ Araiza&Cruz-Conzález(2011) found that there are more aerosolsyear., \citet{Araiza11} found that there are more aerosols.
+ These results suggest that when the sky is clear of clouds the value of ey is related to the atmosphere transparency., These results suggest that when the sky is clear of clouds the value of $\psi_0$ is related to the atmosphere transparency.
+ For completeness we carried out à statistical analysis of the f(elear) for airmass lower than LO., For completeness we carried out a statistical analysis of the f(clear) for airmass lower than 10.
+ First. we obtained a non symmetrical distribution of hourly clear time showing that 78.6 per cent of the time the skv is completely clear at this airmass interval.," First, we obtained a non symmetrical distribution of hourly clear time showing that 78.6 per cent of the time the sky is completely clear at this airmass interval."
+ We calculated. the first. second and third quartile of [(clear) for different perioclicities.," We calculated the first, second and third quartile of f(clear) for different periodicities."
+ We presented. the results for cach month of the four-vear observing period., We presented the results for each month of the four-year observing period.
+ An annual evele is clearly noticeable: large Luetuations are observed mainly during the summer months., An annual cycle is clearly noticeable: large fluctuations are observed mainly during the summer months.
+ Big variations are also present for 2005 in January. February and December but this trend. is not. reproduced in 2006.," Big variations are also present for 2005 in January, February and December but this trend is not reproduced in 2006."
+ The monthly distribution of (clear) was obtained bv combining the data of the same months for the whole period., The monthly distribution of f(clear) was obtained by combining the data of the same months for the whole period.
+ The median of ((clear) is  0.99 between March and Alay. is relatively poor between June and September. with," The median of f(clear) is $\sim$ 0.99 between March and May, is relatively poor between June and September, with"
+We have now introduced the operator S in oue dimoeusiou.,We have now introduced the operator $\cal S$ in one dimension.
+ Next we address how to solve Eq.(," Next we address how to solve Eq.,"
+39).. and then we ecneralize S to higher dimecusions., and then we generalize $\cal S$ to higher dimensions.
+ We prescut our method in this order because the ecucralization to higher dimensious depeuds ou sole details of the solution process., We present our method in this order because the generalization to higher dimensions depends on some details of the solution process.
+ Iu this section we describe how we solve the svsteui of nonlinear equations(39)., In this section we describe how we solve the system of nonlinear equations.
+. Our procedure is completely standard and requires three iugredieuts: À Raphlson iteration to reduce the nonlinear equations to a linear solve at each iteration. an iterative linear solver. aud the preconditioner for the linear solver.," Our procedure is completely standard and requires three ingredients: A Newton-Raphson iteration to reduce the nonlinear equations to a linear solve at each iteration, an iterative linear solver, and the preconditioner for the linear solver."
+ For these three steps sve eiiploy the software package PETSc[?].., For these three steps we employ the software package \cite{petsc-home-page}.
+ We now comicnt oneach of these three stages., We now comment oneach of these three stages.
+ PETSc|?] iuiplomeuts a Newton-Raphsou iiethod with line searches. similar to the method described in [?]..," \cite{petsc-home-page} implements a Newton-Raphson method with line searches, similar to the method described in \cite{NumericalRecipes}."
+" Given a current guess n44 of the solution. a step proceeds as follows: Compute the residual and linearize around the current guess u,44 of tle solution: The Jacobian ο isa ΡΕ« Npp-dimensional matiis. ΟΕ being the uuuber of degrees of freedom."," Given a current guess $\underline{\bf
+ u}_{\mbox{\footnotesize old}}$ of the solution, a Newton-Raphson step proceeds as follows: Compute the residual and linearize $\cal S$ around the current guess $\underline{\bf
+ u}_{\mbox{\footnotesize old}}$ of the solution: The Jacobian ${\cal J}$ is a $N_{DF}\times N_{DF}$ -dimensional matrix, $N_{DF}$ being the number of degrees of freedom."
+ Next compute a correction du bv solving the linear svstei Finally a line-scarch is performed in the direction of δα., Next compute a correction $\delta\underline{\bf u}$ by solving the linear system Finally a line-search is performed in the direction of $\delta\underline{\bf u}$ .
+ Parametrize the new solution by, Parametrize the new solution by
+"extant. any “aperture effects"" on our resulting analysis should be stall (<<o0.02"") compared to the more dominant effects of photometric errors.","extant, any “aperture effects” on our resulting analysis should be small $\lesssim 0.02^m$ ) compared to the more dominant effects of photometric errors."
+ The last step in preparing these data for analysis was to combine tlie optical imagine. imaging. aud tle spectroscopic observations iuto a single dataset for each cluster field.," The last step in preparing these data for analysis was to combine the optical imaging, near-infrared imaging, and the spectroscopic observations into a single dataset for each cluster field."
+ Since the original spectroscopic targets were determined [rom a separate catalog. we first cross-matehed the spectroscopic catalog with tle optical cataloge usingOm a eerowlneOm annulus techuiqueΤ where the augular distance c was determined using the formula where the augles (0= Declination. o= Right Ascension) have been properly converted into raclians. aud the subscript. s. 'efers to specroscopic target aud the subscript. p. refers to photometric object.," Since the original spectroscopic targets were determined from a separate catalog, we first cross-matched the spectroscopic catalog with the optical catalog using a growing annulus technique where the angular distance $\psi$ was determined using the formula where the angles $\theta \equiv$ Declination, $\phi \equiv$ Right Ascension) have been properly converted into radians, and the subscript, $s$, refers to spectroscopic target and the subscript, $p$, refers to photometric object."
+ In all cases. the matei-up occurred for very sinall augles because spectroscopic targets were selected [rom the optical images.," In all cases, the match-up occurred for very small angles because spectroscopic targets were selected from the optical images."
+ The cross-identiication between the infrared aud optical catalogs. however. was slightly more complicated. althoug bit was done in au identical mauner.," The cross-identification between the infrared and optical catalogs, however, was slightly more complicated, although it was done in an identical manner."
+ This was due o the dillerent telescopes ancl detectors used curing the observations. aud the correspoucing geojetric. distortions present between the two catalogs.," This was due to the different telescopes and detectors used during the observations, and the corresponding geometric distortions present between the two catalogs."
+ Furthermore. the variation of tlie spectral eerey distributious of the sources in our sample between the { aud A-f bands cau be quiten large. especially for galaxies at relatively high redshift.," Furthermore, the variation of the spectral energy distributions of the sources in our sample between the $I$ and $K^{'}$ bands can be quite large, especially for galaxies at relatively high redshift."
+ As a result. some objects in each catalog remaied unmatched below the augular limit which is approxinately the confusiou beam width (approxiuately live arcsecouds).," As a result, some objects in each catalog remained unmatched below the angular limit which is approximately the confusion beam width (approximately five arcseconds)."
+ The technique that we have developed [or determining cluster membership requires redshift estimates for all galaxies in the field., The technique that we have developed for determining cluster membership requires redshift estimates for all galaxies in the field.
+ Although the techuique is iudependent of the actual redshift estimation techulque. we have use the empirical photometric redshift teclinique 1999) to generate recIshift estimates for all galaxies in our sample.," Although the technique is independent of the actual redshift estimation technique, we have used the empirical photometric redshift technique \citep{brunner99} to generate redshift estimates for all galaxies in our sample."
+ This is the result of the simplicity oL the empirical techulque in determiniug redshifts (or a photometric sample of £galaxies for which au adequate training set exists., This is the result of the simplicity of the empirical technique in determining redshifts for a photometric sample of galaxies for which an adequate training set exists.
+ Fiuthermore. the empirical techulque. since it is based ouly ou the data uncer investigation. ]s more rost than:dteruative techuiques to uucertalnties iu the spectral euergv distributions of galaxies.," Furthermore, the empirical technique, since it is based only on the data under investigation, is more robust than alternative techniques to uncertainties in the spectral energy distributions of galaxies."
+ These alteriative techniques. such as templiesbased: inetlods. could also be used. as ος as a rea]stic estinate of the error iu the estlunaed redshift cau be determiuec.," These alternative techniques, such as template-based methods, could also be used, as long as a realistic estimate of the error in the estimated redshift can be determined."
+ The empirical plotometric recsilt techique requires a calibration sampe to determine the coellicients of the empirical relatiou., The empirical photometric redshift technique requires a calibration sample to determine the coefficients of the empirical relation.
+ For this work. we have used tlie spectroscopic data which," For this work, we have used the spectroscopic data which"
+"when the trajectory is given as finite time series for the position, velocity and acceleration.","when the trajectory is given as finite time series for the position, velocity and acceleration."
+" The algorithm is based on formulating an instantaneous at each point on the trajectory, and makes use of the concept of the photon formation length to evaluate this quantity."," The algorithm is based on formulating an instantaneous at each point on the trajectory, and makes use of the concept of the photon formation length to evaluate this quantity."
+ It is suitable for use in post-processing the output from a particle-in-cell simulation., It is suitable for use in post-processing the output from a particle-in-cell simulation.
+ Two problems arise with a numerical evaluation of the radiation., Two problems arise with a numerical evaluation of the radiation.
+" At high frequencies, the finite time-resolution of the trajectory is a limitation."," At high frequencies, the finite time-resolution of the trajectory is a limitation."
+" For relativistic particles, this problem can be alleviated by suitable grouping of the terms associated with the slowly and rapidly varying components of the instantaneous power, which improves the stability and accuracy."," For relativistic particles, this problem can be alleviated by suitable grouping of the terms associated with the slowly and rapidly varying components of the instantaneous power, which improves the stability and accuracy."
+" But even with the appropriate grouping of the terms, if the time-resolution of the trajectory cannot be improved indefinitely, a purely numerical evaluation still fails at sufficiently high frequency."," But even with the appropriate grouping of the terms, if the time-resolution of the trajectory cannot be improved indefinitely, a purely numerical evaluation still fails at sufficiently high frequency."
+" Fortunately, it is precisely in this range that the instantaneous power can safely be evaluated using the synchrotron approximation."," Fortunately, it is precisely in this range that the instantaneous power can safely be evaluated using the synchrotron approximation."
+ Subsequent integration of this quantity does not present a difficulty., Subsequent integration of this quantity does not present a difficulty.
+" However, for non-relativistic particles, or for frequencies comparable to the instantaneous angular frequency of the emitting particle, additional terms enter into the expression for the instantaneous power At low frequencies, a (40))limitation is imposed when the trajectory is known only within a finite time interval."," However, for non-relativistic particles, or for frequencies comparable to the instantaneous angular frequency of the emitting particle, additional terms enter into the expression for the instantaneous power \ref{firstorderint}) ) At low frequencies, a limitation is imposed when the trajectory is known only within a finite time interval."
+" This intrinsic restriction cannot be removed: the minimum frequency at which a light curve can be computed is roughly 1052/T, where T is the length of the available time series."," This intrinsic restriction cannot be removed: the minimum frequency at which a light curve can be computed is roughly $10\gamma^2/T$ , where $T$ is the length of the available time series."
+" In current PIC simulations, however, the neglect of the collective response of the plasma to the propagating waves, which leads to effects such as Razin-Tsytovich suppression and transition radiation, is also likely to be important."," In current PIC simulations, however, the neglect of the collective response of the plasma to the propagating waves, which leads to effects such as Razin-Tsytovich suppression and transition radiation, is also likely to be important."
+" These should intervene at frequencies below roughly ~γωρ. where Wp=/Amne?/m is the plasma frequency and n the number density, but to date it does not appear feasible to account for such effects self-consistently."," These should intervene at frequencies below roughly $\sim\gamma\omega_{\rm p}$, where $\omega_{\rm p}=\sqrt{4\pi n e^2/m}$ is the plasma frequency and $n$ the number density, but to date it does not appear feasible to account for such effects self-consistently."
+" For particles of Lorentz factor 100, these estimates imply that the intrinsic restriction is more important than the neglect of collective effects only when time-series shorter than about 10? plasma cycles are used to compute the emitted radiation."," For particles of Lorentz factor $100$, these estimates imply that the intrinsic restriction is more important than the neglect of collective effects only when time-series shorter than about $10^3$ plasma cycles are used to compute the emitted radiation."
+" Our results obtained for 3D magnetostatic turbulence confirm that one observational signature of short length-scale turbulence, is the presence of a high-frequency power-law tail in the mono-energetic emission spectrum Fleishman(2006a)."," Our results obtained for 3D magnetostatic turbulence confirm that one observational signature of short length-scale turbulence, is the presence of a high-frequency power-law tail in the mono-energetic emission spectrum \cite{fleishman06b}."
+". For a power law of electrons dn/dyc4?, one expects a synchrotron power-law spectrum £F,cw?, where s=—1)/2, for frequencies below the roll-over frequency of(p the maximum energy electrons."," For a power law of electrons $\diff n/\diff\gamma \propto \gamma^{-p}$, one expects a synchrotron power-law spectrum $F_{\omega}\propto \omega^{-s}$, where $s=(p-1)/2$, for frequencies below the roll-over frequency of the maximum energy electrons."
+" Observations of GRBs place this index in the range 2«p2.8, corresponding to a spectral index of 0.5«s0.9."," Observations of GRBs place this index in the range $2<p<2.8$, corresponding to a spectral index of $0.5<s<0.9$."
+" Thus, unless the turbulence index is extremely hard, a«1, the photon spectrum will not harden at high frequencies, and this observational signature may be difficult to distinguish from a cut-off."," Thus, unless the turbulence index is extremely hard, $\alpha<1$, the photon spectrum will not harden at high frequencies, and this observational signature may be difficult to distinguish from a cut-off."
+" The synchrotron spectrum of particles radiating in a uniform field is nowhere harder than an w!/? power law, which, since harder spectra have been observed in gamma-ray bursts, has led to the discussion of a synchrotron ‘line of citeppreeceetal98.."," The synchrotron spectrum of particles radiating in a uniform field is nowhere harder than an $\omega^{1/3}$ power law, which, since harder spectra have been observed in gamma-ray bursts, has led to the discussion of a synchrotron line of \\citep{preeceetal98}."
+" Our results confirm that this is generally true for isotropic particle distributions in large scale, static, 3D turbulence."," Our results confirm that this is generally true for isotropic particle distributions in large scale, static, 3D turbulence."
+" However, in agreement with other treatments (Fleishman&Urtiev2010) we find that, in the presence of large amplitude turbulence on short length-scales, the low-frequency asymptote can diverge from this value."," However, in agreement with other treatments \citep{fleishmanurtiev10}
+ we find that, in the presence of large amplitude turbulence on short length-scales, the low-frequency asymptote can diverge from this value."
+ We find low frequencyspectra that are typically harder than w!/®., We find low frequencyspectra that are typically harder than $\omega^{1/3}$.
+ For fields with a>>1 the spectrum exhibits a gradual hardening with the slope increasing approximately 0.05 per decade in frequency., For fields with $a\gg1$ the spectrum exhibits a gradual hardening with the slope increasing approximately $0.05$ per decade in frequency.
+" For fields with strength parameters a=1, low frequency power-law asymptotes are produced, ζωο:uw, with 1/3€q< 1/2."," For fields with strength parameters $a\gtrsim1$, low frequency power-law asymptotes are produced, $F_{\omega}\propto\omega^q$, with $1/3\leq q\leq 1/2$ ."
+" For fields with a<1 slightly larger values of q appear to be possible, although shocks with such small strength parameters are poor accelerators (Kirk&Reville2010)."," For fields with $a<1$ slightly larger values of $q$ appear to be possible, although shocks with such small strength parameters are poor accelerators \citep{kirkreville10}."
+. Spectra as hard as w? are known to be produced by weak(a< turbulence that can be factorized into 2D and 1D components1) (Fleishman 2006)..," Spectra as hard as $\omega^1$ are known to be produced by weak$a\ll1$ ) turbulence that can be factorized into 2D and 1D components \citep{fleishman06,medvedev06}. ."
+" However, they do not arise inour results, which are based on a fully 3D turbulence model."," However, they do not arise inour results, which are based on a fully 3D turbulence model."
+of 3075 Ix is used. equation 16 vields a dust sublimation radius of 4.5011 em and flave-disk interaction becomes probable.,"of 3075 K is used, equation 16 yields a dust sublimation radius of $4.5 \times 10^{11}$ cm and flare-disk interaction becomes probable."
+ The case for flave-cisk interaction in J182016.5-161003 remains ambiguous., The case for flare-disk interaction in J182016.5-161003 remains ambiguous.
+ For the second long flare. COUP 1246 with flare loop length 5.0xLOM. Favata (2005) publish a mass of 0.2 M. and a stellar radius of 1.6 2...," For the second long flare, COUP 1246 with flare loop length $5.0 \times 10^{11}$, Favata (2005) publish a mass of 0.2 $M_{\odot}$ and a stellar radius of 1.6 $R_{\odot}$."
+ Using a corresponding Siess model temperature of 4000 Ix. the corresponding dust destruction radius is 4.0xLot cm. which is easily within reach of the flare.," Using a corresponding Siess model temperature of 4000 K, the corresponding dust destruction radius is $4.0 \times 10^{11}$ cm, which is easily within reach of the flare."
+" Five of the 14 flares in our sample associated with disked YSOs have Lzx0.5/2,. much like the flares described by Pandey Singh (2008) on active evolved stars."," Five of the 14 flares in our sample associated with disked YSOs have $L \le 0.5 R_\star$, much like the flares described by Pandey Singh (2008) on active evolved stars."
+ Compact Lares can be wholly contained by the YSO's dipole field regardless of the presence of cireumstellar accretion disks: for Chis reason compact flares are generally considered analogous to coronal events seen on ZAMS or AIS stus., Compact flares can be wholly contained by the YSO's dipole field regardless of the presence of circumstellar accretion disks; for this reason compact flares are generally considered analogous to coronal events seen on ZAMS or MS stars.
+ The wide range of loop lengths in the sample suggests that either of two types of flares may occur on clisked YSOs: compact flares analogous to flares on evolved stars or long and the result of star-disk magnetic connections., The wide range of loop lengths in the sample suggests that either of two types of flares may occur on disked YSOs: compact flares – analogous to flares on evolved stars – or long and the result of star-disk magnetic connections.
+ Star-disk magnetic coupling can explain the YSO ¢lass-specific distribution of loop lengths seen in Figure 5.., Star-disk magnetic coupling can explain the YSO class-specific distribution of loop lengths seen in Figure \ref{LoopL_hist}.
+ The disks of Class I and ID YSOs have inner radii close to surface of the central star and are thus far more likely to support the extended magnetic structures of star-clisk coupling than the meager or non-existent disks associated will Class IID YSOs., The disks of Class I and II YSOs have inner radii close to surface of the central star and are thus far more likely to support the extended magnetic structures of star-disk coupling than the meager or non-existent disks associated with Class III YSOs.
+ If star-disk magnetic coupling drives long fares we should not see any long Lares associated with Class ΤΗ YSOs., If star-disk magnetic coupling drives long flares we should not see any long flares associated with Class III YSOs.
+ The loop length histograms of Figure 5. are evidence lor (he existence of star-disk magnetic coupling. as [lares longer than 1.5x10! em are only observed on Class ] and II YSOs.," The loop length histograms of Figure \ref{LoopL_hist} are evidence for the existence of star-disk magnetic coupling, as flares longer than $1.5 \times 10^{11}$ cm are only observed on Class I and II YSOs."
+ From these resulis and the results published in (ae COUP survey (855.2). a consistent picture of star-disk magnetic interaction in YSOs emeges.," From these results and the results published in the COUP survey 5.2), a consistent picture of star-disk magnetic interaction in YSOs emeges."
+ As part of normal chromospheric aclivilv on a voung star. the local magnetic field is arranged in arcade loops analagous io those seen on the Sun.," As part of normal chromospheric activity on a young star, the local magnetic field is arranged in arcade loops analagous to those seen on the Sun."
+ Convection on (he surface of the star causes shullling of the loop footprint. which eventually stresses (he magnetic field loops and causes a reconnection event and flare.," Convection on the surface of the star causes shuffling of the loop footprint, which eventually stresses the magnetic field loops and causes a reconnection event and flare."
+ Since YSOs have large surfaces compared to Sun-like MS stars. convective aclivily (ancl shuffling of the arcade loop footprint) is comparatively reduced.," Since YSOs have large surfaces compared to Sun-like MS stars, convective activity (and shuffling of the arcade loop footprint) is comparatively reduced."
+ If the surface underneath the loop footprint is calm enough. and if (μον are equatorially located. the magnetic loops can grow large enough to interact wilh the inner edge of disk by drawing oul ionized material.," If the surface underneath the loop footprint is calm enough, and if they are equatorially located, the magnetic loops can grow large enough to interact with the inner edge of disk by drawing out ionized material."
+ The presence or absence of a circumstellar disk plavs no role in determining high-contrast flare energetics. but should a field line reach out and find disk material. the disk supports this extended magnetic structure and the accompanying extended. high-contrast flare once the flux tube ruptures.," The presence or absence of a circumstellar disk plays no role in determining high-contrast flare energetics, but should a field line reach out and find disk material, the disk supports this extended magnetic structure and the accompanying extended high-contrast flare once the flux tube ruptures."
+ For a detailed discussion of scenarios in which magnetic fieldl loops can grow to several stellar radii. see MacNeice (2004).," For a detailed discussion of scenarios in which magnetic field loops can grow to several stellar radii, see MacNeice (2004)."
+ο The SER of galaxies brighter than LÍ is the same (<0.1 dex) iu D. I. (I1B) and K selected catalogues.,"$\bullet$ The SFR of galaxies brighter than $L_\ast^I$ is the same $\lsim 0.1$ dex) in B, I, (I+B) and K selected catalogues."
+" This indicates that present optical aud NIR survevs are uulikely to have iissed a substantial fraction population of inassive star formüus objects (with the possible exception of heavily dust-cushrouded e The total SER inteerated over all galaxy Iuuiuosities is the same iu the D. 1, aud. (11D) selected catalogues and is lower iu the IE-solected catalogue by 0.2 dex."," This indicates that present optical and NIR surveys are unlikely to have missed a substantial fraction population of massive star forming objects (with the possible exception of heavily dust-enshrouded $\bullet$ The total SFR integrated over all galaxy luminosities is the same in the B, I, and (I+B) selected catalogues and is lower in the K-selected catalogue by 0.2 dex."
+ This difference originates at luminosities lower than L which implies that I-selected. surveys nüss a significant fraction of star-forming lower-luninosity ο At all redshifts. bhnuuinous galaxies (2>L.) contribute only ~ to the total SER. i.c. the integrated SER of L«L. galaxiesi is a factor of —2 higher than the one of L7L. ο Our fits to the ΕΡΕ luninosity functious sugeest a flat füut-eund slope of a=LOFTcx0.01 in contrast to the asstumed slope of à~1.6 in the literature.," This difference originates at luminosities lower than $L_\ast$ which implies that K-selected surveys miss a significant fraction of star-forming lower-luminosity $\bullet$ At all redshifts, luminous galaxies $L>L_\ast$ ) contribute only $\sim \frac{1}{3}$ to the total SFR, i.e. the integrated SFR of $L<L_\ast$ galaxies is a factor of $\sim 2$ higher than the one of $L>L_\ast$ $\bullet$ Our fits to the FDF luminosity functions suggest a flat faint-end slope of $\alpha=-1.07 \pm 0.04$ in contrast to the assumed slope of $\alpha \sim -1.6$ in the literature."
+ This iuplies that past determinations have overestimated the SER by a factor e The SER is approximately constant over the redshift range dl<2<Land drops by about around +=1.5. if dust corrections coustant with redshift are assumed.," This implies that past determinations have overestimated the SFR by a factor $\bullet$ The SFR is approximately constant over the redshift range $1\le z\le4$ and drops by about around $z=4.5$, if dust corrections constant with redshift are assumed."
+" We thauk the anouvious referee for lis helpful cohunents which improved the presentation of the paper considerably,", We thank the anonymous referee for his helpful comments which improved the presentation of the paper considerably.
+ This work was supported by SED 375 of the DFC. by the Cerman Ministry for Science aud Education (BAIBF) through DESY uuder the project OFAE2PDA/S. and by the Deutsche Forschuugsgecimiemschaft σον the project SCIEN 312/31 M. σον) und T. Exbeu thank C. Wolf for providing some of the optical images used here.," This work was supported by SFB 375 of the DFG, by the German Ministry for Science and Education (BMBF) through DESY under the project 05AE2PDA/8, and by the Deutsche Forschungsgemeinschaft under the project SCHN 342/3--1 M. Schirmer und T. Erben thank C. Wolf for providing some of the optical images used here."
+ Observations have been carried out using the Χαν Large Telescope at the ESO Paranal Observatory under Program ID(s): LP16s.A-0[85, Observations have been carried out using the Very Large Telescope at the ESO Paranal Observatory under Program ID(s): LP168.A-0485
+This paper is the third in a series in which we try to unveil the past evolution of the orbits of the giant planets using the current dynamical structure of the Solar System.,This paper is the third in a series in which we try to unveil the past evolution of the orbits of the giant planets using the current dynamical structure of the Solar System.
+" In these works, we do not assume a priori any preferred model other than the well accepted fact that the giant planets were much closer to each other in the past and somehow moved towards their current orbital radii."," In these works, we do not assume a priori any preferred model other than the well accepted fact that the giant planets were much closer to each other in the past and somehow moved towards their current orbital radii."
+" This radial migration of the giant planets is believed to be the last major event that sculpted the structure of the Solar System (Levison et al.,"," This radial migration of the giant planets is believed to be the last major event that sculpted the structure of the Solar System (Levison et al.,"
+" 2001: Gomes et al.,"," 2001; Gomes et al.,"
+ 2005: Strom et al..," 2005; Strom et al.,"
+" 2005). an assumption which is implicit in our In two previous studies, Morbidelli et al. ("," 2005), an assumption which is implicit in our In two previous studies, Morbidelli et al. ("
+2009) and Brasser et al. (,2009) and Brasser et al. (
+"2009), henceforth labelled Paper I and Paper Il respectively, we investigated how the giant planets (Paper D) and the terrestrial planets (Paper ID) could have achieved orbits with their current secular properties i.e. with their current frequencies and amplitudes in eccentricity and inclination.","2009), henceforth labelled Paper I and Paper II respectively, we investigated how the giant planets (Paper I) and the terrestrial planets (Paper II) could have achieved orbits with their current secular properties i.e. with their current frequencies and amplitudes in eccentricity and inclination."
+" The orbits of the planets are not fixed ellipses: they undergo precession, and the eccentricities and inclinations have long-term (secular) oscillations that, to a first approximation, are described by the Lagrange-Laplace theory: Here t is the immaginary unit and ej.%.σι.Qe are, respectively the eccentricity. inclination, longitude of perihelion and longitude of the node of the planet &. g; and s; are the secular"," The orbits of the planets are not fixed ellipses: they undergo precession, and the eccentricities and inclinations have long-term (secular) oscillations that, to a first approximation, are described by the Lagrange-Laplace theory: Here $\iota$ is the immaginary unit and $e_k, i_k, \varpi_k, \Omega_k$ are, respectively the eccentricity, inclination, longitude of perihelion and longitude of the node of the planet $k$, $g_j$ and $s_j$ are the secular"
+intrinsically simpler to measure. since for older ages the red clump is more populated. hence easier to locate precisely.,"intrinsically simpler to measure, since for older ages the red clump is more populated, hence easier to locate precisely."
+" At intermediate ages. like in the case of NCGC2506. there seems instead to be a discrepaney. (see Table""6). with ages from 9 V too high."," At intermediate ages, like in the case of NGC2506, there seems instead to be a discrepancy (see Table 6), with ages from $\delta$ V too high."
+ Tre same problems exists the οςο relation (their eq., The same problems exists for the CC94 relation (their eq.
+ 3): their measured AVW=1.75 for NCGXC250(i vields an age of about 2.5- Civr. which decreases to 1.8 Civr using instead 1.5 (our value).," 3): their measured $\Delta$ V=1.75 for NGC2506 yields an age of about 2.5 Gyr, which decreases to 1.8 Gyr using instead 1.5 (our value)."
+ Comparing ος AV and. ες 98V [or the 26 clusters common to €C€94 and. JPO4. one finds a mean dilference of about 0.15 mag. in the sense of larger values for CC94.," Comparing CC94's $\Delta$ V and JP94's $\delta$ V for the 26 clusters common to CC94 and JP94, one finds a mean difference of about 0.15 mag, in the sense of larger values for CC94."
+ Restricting the comparison to the 6 clusters in common chosen among the 7 best clusters used. by P904 to build the ALAL relation vieles a mean cdilference of esedt mag (excluding the case of NCIO3G680. a scarcely cluster. for which the dillerence is 0.7. mag).," Restricting the comparison to the 6 clusters in common chosen among the 7 best clusters used by JP94 to build the MAI relation yields a mean difference of 0.25 mag (excluding the case of NGC3680, a scarcely populated cluster, for which the difference is 0.7 mag)."
+ For these 6 objects the mean age cillerence is about 0.1 Car., For these 6 objects the mean age difference is about 0.1 Gyr.
+ We then conclude with this careaf about the MAI age ranking system. which could lead to conspicuous errors in individual ages.," We then conclude with this $caveat$ about the MAI age ranking system, which could lead to conspicuous errors in individual ages."
+ Ehe clfort to build a more refined relation is huge. but worth pursuing.," The effort to build a more refined relation is huge, but worth pursuing."
+ ACKNOWLEDCLEAENTS The bulk of the numerical code for CAID simulations 1as been provided by Laura Gregeio., ACKNOWLEDGEMENTS The bulk of the numerical code for CMD simulations has been provided by Laura Greggio.
+ Phe PRANEC. Geneva and. Padova evolutionary tracks were kindly mace available »v their authors.," The FRANEC, Geneva and Padova evolutionary tracks were kindly made available by their authors."
+ The authors would like to thank Dr. ο Reid for a critical reacling of the manuscript ancl Dr. Vincenzo Testa for help with various aspects of the whotometric reductions., The authors would like to thank Dr. Neill Reid for a critical reading of the manuscript and Dr. Vincenzo Testa for help with various aspects of the photometric reductions.
+ We mace use of the digitization of ohetographie data obtained at the VA Schmidt Telescope (Digitized Sky Survey. produced. at. the Space Telescope Science Institute. uncer US Government grant NAG \V-," We made use of the digitization of photographic data obtained at the UK Schmidt Telescope (Digitized Sky Survey, produced at the Space Telescope Science Institute under US Government grant NAG W-21)."
+ ‘This research. has mace use of the Simbacl database. operated at CDS. Strasbourg. France.," This research has made use of the Simbad database, operated at CDS, Strasbourg, France."
+which consistently found. lognormal distributions. Scalo et al. (,"which consistently found lognormal distributions, Scalo et al. ("
+1998). performed one-dimensional simulations of [orced. supersonic. polvtropic turbulence and. uncovered a lognormal PDE for the cases where either the eas was isothermal (4 = 1) or where the Mach. number was small (M 1).,"1998) performed one-dimensional simulations of forced, supersonic, polytropic turbulence and uncovered a lognormal PDF for the cases where either the gas was isothermal $\gamma$ = 1) or where the Mach number was small $\ll$ 1)."
+ Otherwise. when 5.« 1. power laws developed for densities larger than the mean.," Otherwise, when $\gamma <$ 1, power laws developed for densities larger than the mean."
+ Alternatively. Nordlund Padoan (1999) interpreted Sealo ct al," Alternatively, Nordlund Padoan (1999) interpreted Scalo et al."
+os results for the PDs occuring in the ~z 1 case as skewed log-normals and Passot Vázzequez-Semacdeni (1998). provided a mathematical framework for understanding why these distributions arose.,'s results for the PDFs occuring in the $\gamma \neq$ 1 case as skewed log-normals and Passot Vázzquez-Semadeni (1998) provided a mathematical framework for understanding why these distributions arose.
+ Our work extends these investigations on the PDE in the direction. of the cases where the ISM. is constrained, Our work extends these investigations on the PDF in the direction of the cases where the ISM is constrained
+in the central parsec of the Galaxy (Libonate et al.,in the central parsec of the Galaxy (Libonate et al.
+ 1995: Roberts. Yuset-Zacdeh. Coss 1996: IIerbst et al.," 1995; Roberts, Yusef-Zadeh, Goss 1996; Herbst et al."
+ 1993)., 1993).
+ However. IRS 13 has a rather weak οσα line: Roberts et al (1996) thus sugeest that the IT926. source is hotter. denser and more compact than the Bre source (Roberts. Yusef-Zadelh. Coss 1996).," However, IRS 13 has a rather weak $92\alpha$ line; Roberts et al (1996) thus suggest that the $92\alpha$ source is hotter, denser and more compact than the $\gamma$ source (Roberts, Yusef-Zadeh, Goss 1996)."
+ Although its [Fe1T] CLUISSIOML in weak (Stolovy et al., Although its [FeII] emission is weak (Stolovy et al.
+ 1999). IRS 13 is located at the ocal peak of [FcHI cluission (Lutz. Krabbe. Cenzel 1993). suggesting a local ionizing source.," 1999), IRS 13 is located at the local peak of [FeIII] emission (Lutz, Krabbe, Genzel 1993), suggesting a local ionizing source."
+ It should be ioted. however. that for some of the above observations. ine blending (e.g. Πο1-3) with Pa-a aud ITIC12-1) with he 1.GLI ji [Fel] Πιο) can be a significant problem.," It should be noted, however, that for some of the above observations, line blending (e.g. HeI(4-3) with $\alpha$ and HI(12-4) with the 1.644 $\mu$ m [FeII] line) can be a significant problem."
+ From observations of the K-band Wel aud Bre emissiou ines. NOT constructed non-LTE atmospheric inodels of IRS EL aud determined the parameters given in Table 1..," From observations of the K-band HeI and $\gamma$ emission lines, N97 constructed non-LTE atmospheric models of IRS 13E1 and determined the parameters given in Table \ref{tab:naj}."
+ At first elance this mass loss rate is substantially nore than the canonical upper limit of ~1&lo13 vob for WR stars (Nugis 1999)., At first glance this mass loss rate is substantially more than the canonical upper limit of $\sim 1\times10^{-4}~\mathrm{M}_{\sun}$ $^{-1}$ for WR stars (Nugis 1999).
+ Tlowever. a uuuber of arguments exist against such a high estimate.," However, a number of arguments exist against such a high estimate."
+ First N97 did not include wind chunping or lue-blanketine aud as a result probably over estimuted the mass loss rate bv at least a factor of 2 (Morris et al., First N97 did not include wind clumping or line-blanketing and as a result probably over estimated the mass loss rate by at least a factor of 2 (Morris et al.
+ 2000)., 2000).
+ Second. assuming the parameters m Table 1.. the classical stellar wind theory of Wright Barlow (1975) predicts a 13 nuu flux of ~30 ταν. considerably higher than the observed ~10 indy for the combined IRS 13E complex. again sugecsting an overestimate of AL/ex by at least a factor of 2.," Second, assuming the parameters in Table \ref{tab:naj}, the classical stellar wind theory of Wright Barlow (1975) predicts a 13 mm flux of $\sim 30$ mJy, considerably higher than the observed $\sim 10$ mJy for the combined IRS 13E complex, again suggesting an overestimate of $\dot\mathrm{M}/v_\infty$ by at least a factor of 2."
+ Third. if the other CC wind sources have similarly overestimated mass loss rates. lydvodvuaniuecal accretion siunulatious (0.9. Coker Molia. 1997). which predict a arpe Inass accretion rate onto Ser A*. are broueht more iuto agreement with other accretion models (Coker Melia. 2000: Naravan et al.," Third, if the other GC wind sources have similarly overestimated mass loss rates, hydrodynamical accretion simulations (e.g. Coker Melia 1997), which predict a large mass accretion rate onto Sgr A*, are brought more into agreement with other accretion models (Coker Melia, 2000; Narayan et al."
+ 1998). which require a much ower accretion rate.," 1998), which require a much lower accretion rate."
+ For completeness. we also note that if the metallicity of IRS 13E is more than twice solar (sce. e.g.. Shields Ferland 19914: Simpson et al.," For completeness, we also note that if the metallicity of IRS 13E is more than twice solar (see, e.g., Shields Ferland 1994; Simpson et al."
+ 1995). the nasi possible WR inass loss rate could be higher (Maeder Mawvuet 1987).," 1995), the maximum possible WR mass loss rate could be higher (Maeder Maynet 1987)."
+ Tf we asstune the metallicity of the GC is twice solar hen based on the estimated bolometric huuimositv of TRS 13E1 the preseut mass and zero-age mail-sequence CAMS) mass of IRS 13EL are ~60AL: and more han 120 AL.. respectively (Schaller et al 1992).," If we assume the metallicity of the GC is twice solar then based on the estimated bolometric luminosity of IRS 13E1 the present mass and zero-age main-sequence (ZAMS) mass of IRS 13E1 are $\sim 60~\mathrm{M}_{\sun}$ and more than 120 $\mathrm{M}_{\sun}$, respectively (Schaller et al 1992)."
+ However. updated evolutionary models (e.¢.. Mowlavi et al 1998) show that a hieh metallicity star with ZAMS ass of z00Mc will shed so much mass while on the main sequence that it would uever reach the WR. stage.," However, updated evolutionary models (e.g., Mowlavi et al 1998) show that a high metallicity star with ZAMS mass of $\simgt 60~\mathrm{M}_{\sun}$ will shed so much mass while on the main sequence that it would never reach the WR stage."
+ The canonical assumption that the metallicity of the GC is twice solar. however. is based on gas phase abundances.," The canonical assumption that the metallicity of the GC is twice solar, however, is based on gas phase abundances."
+ Carr. Sellereu. Balachandran (2000) further areuc that GC stellar abuudances may in tur be overestimated due to higher stellar rotation rates.," Carr, Sellgren, Balachandran (2000) further argue that GC stellar abundances may in turn be overestimated due to higher stellar rotation rates."
+ For example. high-resolution ucar-intrared spectra of IRS 7. a red supereiaut located in the ceutral parsec. show solar iietallicity.," For example, high-resolution near-infrared spectra of IRS 7, a red supergiant located in the central parsec, show solar metallicity."
+ It is certain that IRS 13E is a luminous object eimibedde in a hot. dense medium. but more than this is not clear.," It is certain that IRS 13E is a luminous object embedded in a hot, dense medium, but more than this is not clear."
+ However. it is possible to explore a few optious which best fit the observations.," However, it is possible to explore a few options which best fit the observations."
+ First. the estimated Iuuinosity ix mass loss rate of IRS 19101 are too laree for a sinele evolve WR stax: they are more like that of an LBV.," First, the estimated luminosity and mass loss rate of IRS 13E1 are too large for a single evolved WR star; they are more like that of an LBV."
+ In addition Libonate et al. (, In addition Libonate et al. (
+1995) noted that the Ik-baud spectrum of IRS 13 resembles that of the LBV. P Cve.,"1995) noted that the K-band spectrum of IRS 13 resembles that of the LBV, P Cyg."
+ Similarly. the FWIIL of the Br-~ line (~350 kin 13 as well as the extent of the cussion region 20.12 pc) (IIerbst et al.," Similarly, the FWHM of the $\gamma$ line $\sim 350$ km $^{-1}$ ) as well as the extent of the emission region $\simlt 0.12$ pc) (Herbst et al."
+ 1993) is consistent with the idea that IRS I3E has receutly nudergone au LBV-like mass ejection. and also tics in well with the radio aud stb-uun observations.," 1993) is consistent with the idea that IRS 13E has recently undergone an LBV-like mass ejection, and also ties in well with the radio and sub-mm observations."
+ Towever. iu spite of this agreement its temperature aud wind velocity are larger than a typical LBV.," However, in spite of this agreement its temperature and wind velocity are larger than a typical LBV."
+ As has been suggestedOO for IRS, As has been suggested for IRS
+wwithChandra.,with.
+ Iu Section ??.. we describe our data aud reduction procedures. as well as several H'chival data sets.," In Section \ref{ana}, we describe our data and reduction procedures, as well as several archival data sets."
+ In Section ?? we analyze the spectra of both the point source and the exteudecd enission. calculate radial profiles. aud examine the extended morphology.," In Section \ref{res} we analyze the spectra of both the point source and the extended emission, calculate radial profiles, and examine the extended morphology."
+ We also describe the :m'operties of a number of trausieut sources that appear in the field., We also describe the properties of a number of transient sources that appear in the field.
+ In Section ?? we discuss our [findings aud their inmiplicatious., In Section \ref{dis} we discuss our findings and their implications.
+ We conclude with a brief sumuinary in Section ??.., We conclude with a brief summary in Section \ref{sum}.
+ We also include Appendix A. in which we calculate the X-ray contribution due to dust scattering in the direction of5039.," We also include Appendix A, in which we calculate the X-ray contribution due to dust scattering in the direction of."
+. Table 1 lists the X-ray observations cousidered iu this work. including five archival observatious (from and Chandra) aud our latest observation.," Table \ref{log} lists the X-ray observations considered in this work, including five archival observations (from and ) and our latest observation."
+ LS 5039 was observed with the Advanced CCD Lnagiug Spectrometer (ACIS) on board on 2009 October 30 (ObsID 106906)., LS 5039 was observed with the Advanced CCD Imaging Spectrometer (ACIS) on board on 2009 October 30 (ObsID 10696).
+ Our observation spauned orbital phases o=0.06[— (using the ephemeris of ?.. where o=0 correspouds to periastrou).," Our observation spanned orbital phases $\phi=0.064$ --0.184 (using the ephemeris of \citealt{2005MNRAS.364..899C}, where $\phi=0$ corresponds to periastron)."
+ The target was observed for kks in timed exposure mode using “very Faint” telemetry format., The target was observed for ks in timed exposure mode using “very faint” telemetry format.
+ To miuinize pile-up ellects. we used a custom 750 pixel wide sub-array and turned. olf all the ACIS chips except I3.," To minimize pile-up effects, we used a custom 750 pixel wide sub-array and turned off all the ACIS chips except I3."
+ The (rame time in this configuration is 2.2110[ss (2.255 exposure time and [ss image transfer time)., The frame time in this configuration is s s exposure time and s image transfer time).
+ The Y-of[set was —2/8. placing the source 3/003 olf axis.," The Y-offset was $-2\farcm8$, placing the source 03 off axis."
+ This resulted in a broader PSF. further helpingOm to mitigatee pile-up.," This resulted in a broader PSF, further helping to mitigate pile-up."
+ There were uo significante backgroundOm flares duriusOm the observation., There were no significant background flares during the observation.
+ The useful scientific exposure time (live-time) was kks., The useful scientific exposure time (live-time) was ks.
+ The data were reduced. and analyzed with the Iuteractive Analysis Observations (CLAO) package 11.2). sine CALDB 1.2.0.," The data were reduced and analyzed with the Interactive Analysis Observations (CIAO) package 4.2), using CALDB 4.2.0."
+ To produce images of the pulsar aud its vicinity at subpixel resolution. we removed the pipeline pixel raudomization aud applied the subpixel resolution tool to split-pixel events in the image (?)..," To produce images of the pulsar and its vicinity at subpixel resolution, we removed the pipeline pixel randomization and applied the subpixel resolution tool to split-pixel events in the image \citep{2001ASPC..251..576M}."
+ The image produced (roi this data is shown as the bottom-right panel in Figure 1.., The image produced from this data is shown as the bottom-right panel in Figure \ref{im}.
+ We downloaded archival observatious of5039.. three by aud two byNXMAI-Newton. see Table 1..," We downloaded archival observations of, three by and two by, see Table \ref{log}."
+ In the case of the data. we used the staudard Level 2 calibrated science event files.," In the case of the data, we used the standard Level 2 calibrated science event files."
+ We used tlie energy range kkeV. which is the optimal range for viewing the extended," We used the energy range keV, which is the optimal range for viewing the extended"
+sPhe Virgo. cluster ollers⋅ the best. opportunity. το study in. detail. large. numbers of. galaxies.. (also chwarl. galaxies). over a small region. of. sky.,The Virgo cluster offers the best opportunity to study in detail large numbers of galaxies (also dwarf galaxies) over a small region of sky.
+ Lt is. the nearest (do17- Alpe. psLikhonov et aL.," It is the nearest $\sim17$ Mpc, Tikhonov et al.,"
+" 2000) cluster. with. several hundreds: of ealaxies. (~1277.. members. Binggeli"". ct al. 1984) and the largest dominant structure of the Local Supercluster."," 2000) cluster with several hundreds of galaxies $\sim 1277$ members, Binggeli et al, 1984) and the largest dominant structure of the Local Supercluster."
+ Lt is an irregularly shaped: cluster. with a high abundance of spiral galaxies among the bright cluster members and a larec population of the total known galaxy number) of ἄννα galaxiese (Dinggeli et al.," It is an irregularly shaped cluster, with a high abundance of spiral galaxies among the bright cluster members and a large population of the total known galaxy number) of dwarf galaxies (Binggeli et al.,"
+ 1985)., 1985).
+ Its size is approximately 10 degrees across the sky. which corresponds to ~3 Alpe.," Its size is approximately 10 degrees across the sky, which corresponds to $\sim 3$ Mpc."
+" The1 crossing time for this cluster is zz0.144, πο et al.."," The crossing time for this cluster is $\approx 0.1
+H_0^{-1}$ (Trentham et al.,"
+ 2002) so that cluster galaxies have had. plenty of opportunities to interact with each other., 2002) so that cluster galaxies have had plenty of opportunities to interact with each other.
+ The huge mass of the cluster (AZ=L2«101A. Fouque. 2001) accelerates 1e member galaxies to very high peculiar velocities. so that -- exhibits the highest. bluc-shift. measured for any galaxy (103258. approaching us at 1600 km/sec).," The huge mass of the cluster $M=1.2 \times 10^{15} M_{\odot}$, Fouque, 2001) accelerates the member galaxies to very high peculiar velocities, so that it exhibits the highest blue-shift measured for any galaxy (IC3258, approaching us at 1600 km/sec)."
+ As suggested [rom its irregular structure. the cluster is made up of at least three subclusters dominated by the bright ellipticals AIST. M86. ΛΗΛ) and thus it ..is probably not in. a state of. equilibrium:eu. it'sD. avery complex unrelaxed system where the central part is. dynamically. old. while. the large surrounding. region ..is not virialized.EN but still ⊲⊲⋅⊀infalling (Bohringer.⋝. 1995).," As suggested from its irregular structure, the cluster is made up of at least three subclusters dominated by the bright ellipticals M87, M86, M49 and thus it is probably not in a state of equilibrium: it's a very complex unrelaxed system where the central part is dynamically old, while the large surrounding region is not virialized, but still infalling (Bohringer, 1995)."
+ The first. observations of the cluster. date. back to Alechain and Messier (ate 18/5. century). who noticed a arge concentration of nebulac in the northern wing of the Virgo constellation.," The first observations of the cluster date back to Mechain and Messier (late $18^th$ century), who noticed a large concentration of nebulae in the northern wing of the Virgo constellation."
+ Phe identification of the cluster as a self-eravitating svstem consisting of hundreds. of galaxies ollowecl soon. after. Hubble's. 1923. discovery. of Cepheids in M31., The identification of the cluster as a self-gravitating system consisting of hundreds of galaxies followed soon after Hubble's 1923 discovery of Cepheids in M31.
+ The first systematic investigation of the cluster was carried out by Shapley and Ames (1932)., The first systematic investigation of the cluster was carried out by Shapley and Ames (1932).
+ Since then he Virgo cluster has been of primary importance for extragalactic astronomy., Since then the Virgo cluster has been of primary importance for extragalactic astronomy.
+ Lt has been subject to many studies at optical. radio. LR. and X-ray wavelengths aimed at aclelressing its galaxy population. evolution and gas content (for the most recent studies see Gavazzi οἱ al. 2002 (optical multiA): Van Driel et al..," It has been subject to many studies at optical, radio, IR, and X-ray wavelengths aimed at addressing its galaxy population, evolution and gas content (for the most recent studies see Gavazzi et al, 2002 (optical $\lambda$ ); Van Driel et al.,"
+ 2000 (radio) Tulls οἱ al..," 2000 (radio) Tuffs et al.,"
+ 2002 (I1): Shibata et al. 2001 (NX-rav)).," 2002 (IR); Shibata et al, 2001 (X-ray))."
+ The most complete optical survey of the Virgo cluster, The most complete optical survey of the Virgo cluster
+ight from a compact source.,light from a compact source.
+ We call this region he NE miror hereafter in this paper., We call this region the NE mirror hereafter in this paper.
+ However. in other regions. the polarization oteru is verv usual.," However, in other regions, the polarization pattern is very unusual."
+ At about 0.75 west from he hotspot. there is a region with a “racial” ottern of large polarization degrees.," At about $0.''5$ west from the hotspot, there is a region with a “radial” pattern of large polarization degrees."
+ We call this a lavec-P radial feature. for couvenicuce in this oxvper.," We call this a large-P radial feature, for convenience in this paper."
+ This is primarily due to the radialb-shape counts in the POLI20 polarizer images at this yosition (counts at this position are elongated along the radial direction with respect to the xiehtest center) where the count rate is ~0.2% of he peak count rate., This is primarily due to the radial-shape counts in the POL120 polarizer images at this position (counts at this position are elongated along the radial direction with respect to the brightest center) where the count rate is $\sim 0.2$ of the peak count rate.
+ A similar but weaker feature is also seen in the POLGO images., A similar but weaker feature is also seen in the POL60 images.
+ This feature is seen du both of the E312W and F275W (but weaker) filter images. so if is not considered to be due to iucideutal high backeround couuts at this position.," This feature is seen in both of the F342W and F275W (but weaker) filter images, so it is not considered to be due to incidental high background counts at this position."
+ The ceutral 0.71 scale region consists of a few cistinet features. as discussed by Heckiiau et al. (," The central $0.''4$ scale region consists of a few distinct features, as discussed by Heckman et al. ("
+1997): (1) a bright poiutlike source. which we call here a hotspot. (2) —0.72 scale fuzz (or brokeu ving feature as Deckiman et al.,"1997): (1) a bright pointlike source, which we call here a hotspot, (2) $\sim 0.''2$ scale fuzz (or broken ring feature as Heckman et al."
+" called it) clongated roughly north-south around the hotspot. (3) a fainter spiral arma of ~(0,71 scale."," called it) elongated roughly north-south around the hotspot, (3) a fainter spiral arm of $\sim
+0.''4$ scale."
+ This central reeion also has a somewhat raclial-like PA pattern., This central region also has a somewhat radial-like PA pattern.
+ We show shehltly zoomed polarization maps iu Figures 3. and 1L. where the polarized fiux in each iu ds drawn aud the regions of statistical S/N in P down to Ἰ are included to show the detailed »olarizatfion measurement at this central region.," We show slightly zoomed polarization maps in Figures \ref{fig_f342_papf} and \ref{fig_f275_papf}, where the polarized flux in each bin is drawn and the regions of statistical S/N in $P$ down to 1 are included to show the detailed polarization measurement at this central region."
+" Iu the following. we infer that the radial-like PA vattern in the central 0.""1 reeion is primarily due o an imstrinenutal effect from the compact bright ioftspot | fuzz feature. and also infer that the arec-P radial feature is probably due to the same effect."," In the following, we infer that the radial-like PA pattern in the central $0.''4$ region is primarily due to an instrumental effect from the compact bright hotspot + fuzz feature, and also infer that the large-P radial feature is probably due to the same effect."
+ The FOC polarimetry is implemented by taking nuages through three different polarizers. called POLO. POL60. aud POLI20. which have axes of three different directions with respect to the detector. Cor 18073). (or 2107)). and 1207(or 3007)) counter-clockwise from the image τμdirection.," The FOC polarimetry is implemented by taking images through three different polarizers, called POL0, POL60, and POL120, which have axes of three different directions with respect to the detector, (or ), (or ), and (or ) counter-clockwise from the image $+x$ -direction."
+ Figures σι are all preserved to have the detector coordinate direction (so the north isnof to the top)., Figures \ref{fig_f342_papol}$ $\sim$ \ref{fig_f275_papf} are all preserved to have the detector coordinate direction (so the north is to the top).
+" The nuages through the three polarizers are known to have differeut poiut PAxead functions (PSES). which results in viclding ""jiustrumenutal polarization” when the polarization of a compact source is measured with small apertures (lodge 1995)."," The images through the three polarizers are known to have different point spread functions (PSFs), which results in yielding “instrumental polarization” when the polarization of a compact source is measured with small apertures (Hodge 1995)."
+ Iu fact. this difference in the PSFs would also viel) some iustrunueutal features when polarization is calculated at the reeions around a bright source.," In fact, this difference in the PSFs would also yield some instrumental features when polarization is calculated at the regions around a bright source."
+ This iustruiieutal effect has not been discussed in the The PSFs through POLGO aud POLI20 were found to be slightly elougated along the polarizer axis (for nuages of one star. see Fig.l iu Ποσο 1995). thoueh for POLGO the structure in the wing appears to be more complicated.," This instrumental effect has not been discussed in the The PSFs through POL60 and POL120 were found to be slightly elongated along the polarizer axis (for images of one star, see Fig.1 in Hodge 1995), though for POL60 the structure in the wing appears to be more complicated."
+ This csscutially seenis to result in having an iustrünental raclia PA pattern in the PSF wines., This essentially seems to result in having an instrumental radial PA pattern in the PSF wings.
+" This is shown in Figues 5 and 6 with the “iustrmucuta polarization and ""iunstruüneutal polarized flux”. respectively. when the polarizations are calculate at the PSF wines (this is au artifact that emerges from calculation. aud not the plivsical polarization)."," This is shown in Figures \ref{fig_psf_papol} and \ref{fig_psf_papf} with the “instrumental polarization” and “instrumental polarized flux”, respectively, when the polarizations are calculated at the PSF wings (this is an artifact that emerges from calculation, and not the physical polarization)."
+ The iustrunenutal polarized flux is very nal compared to the total flux of the source., The instrumental polarized flux is very small compared to the total flux of the source.
+ Figure G shows this ratio when binued with 10 pixels (οπι SOLU. ," Figure \ref{fig_psf_papf}
+ shows this ratio when binned with 10 pixels $0.''14 \times
+0.''14$ )."
+The ratio is at about the 1.ο level at the bins adjaceut to the peak bin., The ratio is at about the $1-2$ level at the bins adjacent to the peak bin.
+ When the source is more extended and the intensity eracdieunt is mach sinaller. this iustrauental effect will be more smeared aud become smaller.," When the source is more extended and the intensity gradient is much smaller, this instrumental effect will be more smeared and become smaller."
+ For Mrk 177. however. the intensity geradieut around the hotspot | fuzz secs to be large chough that we see this effect.," For Mrk 477, however, the intensity gradient around the hotspot + fuzz seems to be large enough that we see this effect."
+ The PA pattern seen in the central region (Figs., The PA pattern seen in the central region (Figs.
+ 1 aud 2)) actually resembles the PA pattern in the PSF wines (Fig. 5)., \ref{fig_f342_papol} and \ref{fig_f275_papol}) ) actually resembles the PA pattern in the PSF wings (Fig. \ref{fig_psf_papol}) ).
+ Note that they are all in the same detector coordinate directions aud with the same size bins. so we can directly compare them.," Note that they are all in the same detector coordinate directions and with the same size bins, so we can directly compare them."
+ These features iu the PSF wings would depend on the telescope focus/breathing. but the overall correspondence suggests that our assessment of the PSF effect is," These features in the PSF wings would depend on the telescope focus/breathing, but the overall correspondence suggests that our assessment of the PSF effect is"
+ Iu this paper we report on the serendipitous discovery of a spiral structure in the Virgo cluster dawarf elliptical IC3328., In this paper we report on the serendipitous discovery of a spiral structure in the Virgo cluster dwarf elliptical IC3328.
+ The presence of spiral) structure provides colmpclling evidence for the disk nature of that particular dwart galaxy., The presence of spiral structure provides compelling evidence for the disk nature of that particular dwarf galaxy.
+ The observatious are described aud the light distribution is analysed in Sect.2., The observations are described and the light distribution is analysed in Sect.2.
+ Iu Sect.3 we estimate the kinematical properties of the dwarf galaxy. from the observed light distribution. assuming a likely value for the masse-to-lieht ratio and distance.," In Sect.3 we estimate the kinematical properties of the dwarf galaxy from the observed light distribution, assuming a likely value for the mass-to-light ratio and distance."
+ We speculate on the origin of the spiral pattern iu Sect.l., We speculate on the origin of the spiral pattern in Sect.4.
+ The concludiug section deals with the ranifications for the dwarf elliptical taxonomy arising from IC3328., The concluding section deals with the ramifications for the dwarf elliptical taxonomy arising from IC3328.
+ IC3328 is known as an carly-type dwarf galaxy in the Virgo cluster. morphologically classified as dELT.N (Vireo Cluster Catalog 856. Bingecli et al.," IC3328 is known as an early-type dwarf galaxy in the Virgo cluster, morphologically classified as dE1,N (Virgo Cluster Catalog 856, Binggeli et al."
+ 1985)., 1985).
+ The redshift of ον=972432 loqnade this galaxy a probable cluster member aud a good candidate for a more refined distance determunation based ou the Surface Brightness Fluctuations (SBF) method., The redshift of $v_\odot=972\pm 32$ made this galaxy a probable cluster member and a good candidate for a more refined distance determination based on the Surface Brightness Fluctuations (SBF) method.
+ Candidates for the SBF method should have smooth and svuuuectric light distributions. and show no obvious dust or star-formine roeglous.," Candidates for the SBF method should have smooth and symmetric light distributions, and show no obvious dust or star-forming regions."
+" The firs step in applving the SBF incthod is to determune the mean 2-D surface xiehtuess distribution by suitable averaging aud then subtracting it from the ealaxyimage. leaving just the fluctuating part which arises from the Poisson distribution of he stellar sources. IC,"," The first step in applying the SBF method is to determine the mean 2-D surface brightness distribution by suitable averaging and then subtracting it from the galaxy image, leaving just the fluctuating part which arises from the Poisson distribution of the stellar sources, ie."
+ the fluctuations due to unresolved stars iu the galaxy., the fluctuations due to unresolved stars in the galaxy.
+ However in the case of IC3328 he residuals were most surprising., However in the case of IC3328 the residuals were most surprising.
+ The composite of three. 100 sec αι images of IC3328 obtained at the VET. with the FORSI 1nulti-mode instrument in Service Mode ou July 13. 1999 under excelleut seeing conditions (0.67) is shown in the left paucl of Fie.l..," The composite of three 400 sec $R$ -band images of IC3328 obtained at the VLT with the FORS1 multi-mode instrument in Service Mode on July 13, 1999 under excellent seeing conditions $0.6\arcsec$ ) is shown in the left panel of \ref{fig1}."
+ The surface brightuess “fluctuations” obtained after subtracting the mean is shown iu the right panel., The surface brightness “fluctuations” obtained after subtracting the mean is shown in the right panel.
+ Clearly “fluctuations” is not the right word to describe +1e exceptionally regular spiral whose amplitude is so low +lat it remained invisible in the orginal exposure., Clearly “fluctuations” is not the right word to describe the exceptionally regular spiral whose amplitude is so low that it remained invisible in the original exposure.
+" The 4.oral structure is confined to the inner 30"" of the ealaxy which has an dsophotal radius of rgoz—807."," The spiral structure is confined to the inner $30\arcsec$ of the galaxy which has an isophotal radius of $r_{R, 27}$."
+ It appears that we have discovered the real nature of IC3328: a seenminuelv dust-free disk galaxy with unusually fait spiral structure., It appears that we have discovered the real nature of IC3328: a seemingly dust-free disk galaxy with unusually faint spiral structure.
+ The initial aualvsis of the light distribution iu [C3328 followed the traditional route. based on stiudard IBAE," The initial analysis of the light distribution in IC3328 followed the traditional route, based on standard IRAF"
+while matter Hllowing backwards attains negative angular momentum.,while matter flowing backwards attains negative angular momentum.
+ Finally. note that because of the narrower stream in the post-impact region of the more rapidly rotating core. stream material mixes more ellicientlv with the core material (Fig.," Finally, note that because of the narrower stream in the post-impact region of the more rapidly rotating core, stream material mixes more efficiently with the core material (Fig."
+ 5 and 6)).," \ref{fig:core}
+ and \ref{fig:core2}) )."
+ This may have been important. consequences. for the nucleosvnthesis in this region (see 5.1 and ID)., This may have been important consequences for the nucleosynthesis in this region (see 5.4 and IP).
+ One of the objectives of our study. is to investigate the nucleosvnthesis in the stream material after it has been heated bv its impact with the core of the primary., One of the objectives of our study is to investigate the nucleosynthesis in the stream material after it has been heated by its impact with the core of the primary.
+ The associated heating by nuclear burning itself may then allect the hydrodynamies of the stream., The associated heating by nuclear burning itself may then affect the hydrodynamics of the stream.
+ As can already be seen from the simulations presented so far. the characteristic time-scale for the stream infallis of order 1000 where the stream material typically spends no more than s.50 100s near the impact region before it bounces or forms a ring orbiting the core (in the case of a fast rotating core).," As can already be seen from the simulations presented so far, the characteristic time-scale for the stream infallis of order $1000\,$ s, where the stream material typically spends no more than $50-100\,$ s near the impact region before it bounces or forms a ring orbiting the core (in the case of a fast rotating core)."
+ Nuclear burning can only allect the stream motion if the characteristic time-scale for nuclear burning is of order or shorter than the stream dynamical time-scale., Nuclear burning can only affect the stream motion if the characteristic time-scale for nuclear burning is of order or shorter than the stream dynamical time-scale.
+ For the largest part of the stream trajectory. the stream material is relatively cold. (for the range of stream models considered).," For the largest part of the stream trajectory, the stream material is relatively cold (for the range of stream models considered)."
+ Nuclear burning only starts to increase dramatically after it has been heated to a temperature comparable to the ambient temperature. which occurs suddenly in the region of impact.," Nuclear burning only starts to increase dramatically after it has been heated to a temperature comparable to the ambient temperature, which occurs suddenly in the region of impact."
+ However. since the characteristic nuclear scale is generally. longer than the local dynamical time-scale. stream. material will have already moved. far away rom the region of impact and hence have little ellect on the slreani dynamics (in the pre-impact region) and the stream rajectorv.," However, since the characteristic nuclear time-scale is generally longer than the local dynamical time-scale, stream material will have already moved far away from the region of impact and hence have little effect on the stream dynamics (in the pre-impact region) and the stream trajectory."
+ Ht mainly. alfects the properties of the ambient matter (its composition and temperature) ancl hence the xickerouncd conditions for the stream on a time-scale that is ong compared to the stream infall time-scale (these effects and the associated. nucleosvnthesis can be followed: with a standard stellar evolution code: see LPS for a description of he nucleosvnthesis calculations), It mainly affects the properties of the ambient matter (its composition and temperature) and hence the background conditions for the stream on a time-scale that is long compared to the stream infall time-scale (these effects and the associated nucleosynthesis can be followed with a standard stellar evolution code; see IPS for a description of the nucleosynthesis calculations).
+ The situation can be different when the stream impacts he core more-or-less racially ancl when the core has a very steep density. &radient. causing a very strong shock and hence very high post-shock temperatures.," The situation can be different when the stream impacts the core more-or-less radially and when the core has a very steep density gradient, causing a very strong shock and hence very high post-shock temperatures."
+ This can cause a nuclear runaway as illustrated in Figure 7.., This can cause a nuclear runaway as illustrated in Figure \ref{fig:nuc}.
+ As the middle panels in this figure show. it triggers a detonation wave that propagates up into the stream. almost completely disrupting it in the process.," As the middle panels in this figure show, it triggers a detonation wave that propagates up into the stream, almost completely disrupting it in the process."
+ Phe expansion wave propagating into the ambient. medium. causes a dramatic expansion. of the envelope material., The expansion wave propagating into the ambient medium causes a dramatic expansion of the envelope material.
+ Once all of the matter burning in this runaway has been consumed. one may expect that the stream re-establishes itself. ancl that the process. repeats. leading to pulsed nuclear burning.," Once all of the matter burning in this runaway has been consumed, one may expect that the stream re-establishes itself and that the process repeats, leading to pulsed nuclear burning."
+ In the tcealizecl case of a stream of constant entropy. the depth of penetration can casily be estimated. from. the balancing of the stream ram pressure with the ambient oessure.," In the idealized case of a stream of constant entropy, the depth of penetration can easily be estimated from the balancing of the stream ram pressure with the ambient pressure."
+ Ln our case. the pressure of the stream material comes. comparable to the ambient pressure well. before he stream reaches its deepest point.," In our case, the pressure of the stream material becomes comparable to the ambient pressure well before the stream reaches its deepest point."
+ 1n this case ijt is etter to use an estimate based on the conservation of the Bernoulli integral for matter along a streamline (though roth estimates would. give simular results)., In this case it is better to use an estimate based on the conservation of the Bernoulli integral for matter along a streamline (though both estimates would give similar results).
+ However. since he stream entropy changes because of the compression by shock waves propagating into the stream. from. the sides. one has to use a mocified Bernoulli integral to estimate the »netration depth.," However, since the stream entropy changes because of the compression by shock waves propagating into the stream from the sides, one has to use a modified Bernoulli integral to estimate the penetration depth."
+ This is possible as long as the stream is wide enough that it is not stronely allected by the 9-1 instability and that core rotation plavs a minor role., This is possible as long as the stream is wide enough that it is not strongly affected by the K-H instability and that core rotation plays a minor role.
+ For this purpose. we rewrite the Bernoulli integral for stream material in the form: where is the velocity of stream material. e the sonic speed of the stream material and 9 the potential.," For this purpose, we rewrite the Bernoulli integral for stream material in the form: + + = + +, where $u$ is the velocity of stream material, $c$ the sonic speed of the stream material and $\Phi$ the potential."
+ The subscript S corresponds to matter in the stream. the subscript TE refers to the initial conditions and. the subseript B to the conditions at the point of deepest penetration.," The subscript S corresponds to matter in the stream, the subscript T refers to the initial conditions and the subscript B to the conditions at the point of deepest penetration."
+ The point of deepest penetration can then be determined by the condition that the pressure of the stream is equal to the pressure in the ambient matter and that the final stream. velocity is zero (consistent with the results of our hyvdrodynamical calculations)., The point of deepest penetration can then be determined by the condition that the pressure of the stream is equal to the pressure in the ambient matter and that the final stream velocity is zero (consistent with the results of our hydrodynamical calculations).
+ For thestellar structure parametrized by the pressure eracdient the corresponding potential can be written as vn) a= where the subscript A corresponds to the ambient matter at the initial point of the stream.," For thestellar structure parametrized by the pressure gradient the corresponding potential can be written as (r) = =, where the subscript $\mbox A$ corresponds to the ambient matter at the initial point of the stream."
+" Assuming that the entropy of the stream changes according to E (4 )yua one can rewrite the Bernoulli integral to find A, as a function of relative. penetration. depth. &=rfrp. and the initial conditions for the stream and ambicnt matter: Llere py is the density of the ambient material at the initial radius of the steam. and 3 is the ratio of gas pressure to total pressure for the ambient matter."," Assuming that the entropy of the stream changes according to ( ) = ( ), one can rewrite the Bernoulli integral to find $K_{\rm s}$ as a function of relative penetration depth, $\delta = r/r_{\rm T}$, and the initial conditions for the stream and ambient matter: Here $\rho_{\mbt A }$ is the density of the ambient material at the initial radius of the steam, and $\beta$ is the ratio of gas pressure to total pressure for the ambient matter."
+" 'To determine the entropy change factor. AN, we performed sets of hydrodsnamical simulations for two parametrized structures of the ambient medium. (apio)=(5.2.1.3) and (4.2:0.9). representing typical profiles in massive CI systems."," To determine the entropy change factor, $K_{\rm s }$, we performed sets of hydrodynamical simulations for two parametrized structures of the ambient medium, $(\alpha_{\rm P };\alpha_{\rm T })=(5.2.1.3)$ and $(4.2;0.9)$ , representing typical profiles in massive CE systems."
+" We then calculated the stream: penetration depth for 3 values of density ratios. n,= 2. 10 and 30. and [or"," We then calculated the stream penetration depth for 3 values of density ratios, $\eta_\rho = 2$ , 10 and 30, and for"
+cddata reduction and analvsis: in 3 (he results of (he spectral analvsis $44 the photoionization modeling of (he spectrum: Discussion and Conclusions follow in 5 and 6.,data reduction and analysis; in 3 the results of the spectral analysis 4 the photoionization modeling of the spectrum; Discussion and Conclusions follow in 5 and 6.
+" Throughout this paper. a concordance cosmology withHj=71 ! |. Q4—0.73. and Q,,=0.27 (Spergeletal.2003) is adopted."," Throughout this paper, a concordance cosmology with$_0=71$ $^{-1}$ $^{-1}$ , $\Omega_{\Lambda}$ =0.73, and $\Omega_m$ =0.27 \citep{spergel03} is adopted."
+ Errors are quoted to confidence for 1 parameter of interest. (i.e. Ay? or AC’= 2.71)., Errors are quoted to confidence for 1 parameter of interest (i.e. $\Delta \chi^{2}$ or $\Delta C=2.71$ ).
+ All the spectral parameters in subsequent sections are «quoted in the rest.[rane of 4445 ἐς=0.0562) unless otherwise state., All the spectral parameters in subsequent sections are quoted in the rest–frame of 445 $z=0.0562$ ) unless otherwise stated.
+ Chandra observed 4445 with the LETG (Low Energy Transmission Grating) for a net exposure of kks between 25 September 3 October 2009. with the ACIS-S detector in the focal plane.," Chandra observed 445 with the LETG (Low Energy Transmission Grating) for a net exposure of ks between 25 September – 3 October 2009, with the ACIS-S detector in the focal plane."
+ The zeroth order image of 4445 at the aim-point of ACIS-S appears unresolved., The zeroth order image of 445 at the aim-point of ACIS-S appears unresolved.
+ No significant source variability was found during the observations. so the data was used.," No significant source variability was found during the observations, so the time--averaged data was used."
+ The +1 order spectra were summed for the LETG respectively. along with their response files.," The $\pm1$ order spectra were summed for the LETG respectively, along with their response files."
+ The resultant. summed (background subtracted) first order count rate for the LETG is 0.0318+0.0004 | over the energy range kkeV. Initially. in order to parameterize (he LETG spectrum of 4445 with simple continuum models. we binned (he spectra to a minimum of LO source counts per bin and apply the V statistic.," The resultant summed (background subtracted) first order count rate for the LETG is $0.0318\pm0.0004$ $^{-1}$ over the energy range keV. Initially, in order to parameterize the LETG spectrum of 445 with simple continuum models, we binned the spectra to a minimum of 10 source counts per bin and apply the $\chi^{2}$ statistic."
+" The contimuun of the LETG spectrum from 4445 can be verv erudely parameterized, bv a broken.powerlaw moclel. as is shown in Figure 1."," The continuum of the LETG spectrum from 445 can be very crudely parameterized by a broken–powerlaw model, as is shown in Figure 1."
+ The spectrum at higher energies. above a break energy of ~1.4 kkeV. (~9 AA)) is verv hard. rising with energv (decreasing waveleneth) with a photon index of [D=—0.64+0.03.," The spectrum at higher energies, above a break energy of $\sim 1.4$ keV $\sim9$ ) is very hard, rising with energy (decreasing wavelength) with a photon index of $\Gamma=-0.64\pm0.03$."
+ Below the break. (he X-ray spectrum is solter. with D—1.13x0.21.," Below the break, the X-ray spectrum is softer, with $\Gamma=1.73\pm0.21$."
+" This initial parameterization of (he data provided a poor fit (with 47/dof=568.5/181. null hypothesis probability 2=1.3x10.Ἱ7. where dof is the number of deerees of Ireedom in the fit),"," This initial parameterization of the data provided a poor fit (with $\chi^{2}/{\rm dof}=568.5/181$, null hypothesis probability $P=1.3\times 10^{-41}$, where dof is the number of degrees of freedom in the fit)."
+ At soft N-rav energies. there is statistically significant scatter around the continuum due to the presence of Likely emission lines. while at higher energies. residuals are present in the data. most notably. emission and absorption in the iron Ix band between kkeV (~2 AA)).," At soft X-ray energies, there is statistically significant scatter around the continuum due to the presence of likely emission lines, while at higher energies, residuals are present in the data, most notably emission and absorption in the iron K band between keV $\sim 2$ )."
+" The observed flux of 4445 from kkeV is S.8S4x10P ας ο, ", The observed flux of 445 from keV is $8.84\times10^{-12}$ $^{-2}$ $^{-1}$ .
+In order to providea more physical representation of the continuum. we fit the spectrum wilh a photoelectric absorption model of the form £(I7)=wabsx(zwabspowl+ pow2).," In order to providea more physical representation of the continuum, we fit the spectrum with a photoelectric absorption model of the form $F(E) = {\rm wabs} \times ({\rm zwabs}\times{\rm pow1} + {\rm pow2})$ ,"
+what Kuhn et al. (,what Kuhn et al. (
+"1998) found in the P4(cos6) component of the distortion from sphericity of the shape of the Sun’s surface; the difference could be in part a product of large errors in the distortion observed, but it is likely to be due mainly to the direct effect of the centrifugal force, and also the magnetic field and accompanying velocity and temperature variations associated with solar activity in the low-density surface layers, which influence the figure but hardly modify the gravitational field, even though the kernels for the high-order multipole moments are more concentrated near the surface.","1998) found in the $P_4(\cos\theta)$ component of the distortion from sphericity of the shape of the Sun's surface; the difference could be in part a product of large errors in the distortion observed, but it is likely to be due mainly to the direct effect of the centrifugal force, and also the magnetic field and accompanying velocity and temperature variations associated with solar activity in the low-density surface layers, which influence the figure but hardly modify the gravitational field, even though the kernels for the high-order multipole moments are more concentrated near the surface."
+ The mean value of J4 in Table | is close to that found by Roxburgh (2001) and Mecheri et al. (, The mean value of $J_4$ in Table 1 is close to that found by Roxburgh (2001) and Mecheri et al. (
+"2004), who used rough approximations to the seismically inferred rotation rate.","2004), who used rough approximations to the seismically inferred rotation rate."
+" The variations in Jg that we have inferred from the GONG and MDI data do not agree: whereas the MDI results suggest a clear variation which is correlated with solar activity, the GONG results reveal none."," The variations in $J_6$ that we have inferred from the GONG and MDI data do not agree: whereas the MDI results suggest a clear variation which is correlated with solar activity, the GONG results reveal none."
+" The mean values of Jg differ by a factor 3, although they are at least of the same order as that found by Roxburgh (2001)."," The mean values of $J_6$ differ by a factor 3, although they are at least of the same order as that found by Roxburgh (2001)."
+" The variation in Jg is anticorrelated with solar activity, and now the variations inferred from both GONG and MDI data look similar."," The variation in $J_8$ is anticorrelated with solar activity, and now the variations inferred from both GONG and MDI data look similar."
+" However, the mean values do not even have the same sign."," However, the mean values do not even have the same sign."
+ With Jig and Jj» there is good agreement between GONG and MDI., With $J_{10}$ and $J_{12}$ there is good agreement between GONG and MDI.
+" The magnitudes of Jg and higher-order moments are comparable with their temporal variations, although we hasten to add that they are comparable also with the estimated 1-c errors in the individual estimates."," The magnitudes of $J_8$ and higher-order moments are comparable with their temporal variations, although we hasten to add that they are comparable also with the estimated $\sigma$ errors in the individual estimates."
+ Therefore they may not be significant., Therefore they may not be significant.
+ But we have presented them because they do tend to show definite temporal trends which are similar in both the GONG and MDI inferences., But we have presented them because they do tend to show definite temporal trends which are similar in both the GONG and MDI inferences.
+" Higher-order multipole moments also exhibit variation, but we do not present them because their reliability is certainly questionable."," Higher-order multipole moments also exhibit variation, but we do not present them because their reliability is certainly questionable."
+" We point out that the contribution from rotation to each multipole moment can be expressed as a bilinear function of the odd degeneracy splitting coefficients of the seismic modes, a function which depends also on the density p(r)."," We point out that the contribution from rotation to each multipole moment can be expressed as a bilinear function of the odd degeneracy splitting coefficients of the seismic modes, a function which depends also on the density $\rho(r)$."
+" Therefore, because the density of the same reference solar model was used to compute all the moments, whether the seismic data used were from GONG or MDI, all the differences amongst the inferences must have their origin solely in the differences amongst the measured splitting coefficients."," Therefore, because the density of the same reference solar model was used to compute all the moments, whether the seismic data used were from GONG or MDI, all the differences amongst the inferences must have their origin solely in the differences amongst the measured splitting coefficients."
+ We have studied the temporal variation in the rotational kinetic energy and angular momentum of the Sun in different regions of the convection zone., We have studied the temporal variation in the rotational kinetic energy and angular momentum of the Sun in different regions of the convection zone.
+ We find that in the outer layers (r2% O0.9Ro) the variation is well correlated with solar activity., We find that in the outer layers $r\ga 0.9R_\odot$ ) the variation is well correlated with solar activity.
+" At high latitudes this correlation continues throughout the convection zone, while at low latitudes (and also when integrated over all latitudes), there is a transition to anticorrelation around r=0.9Ro."," At high latitudes this correlation continues throughout the convection zone, while at low latitudes (and also when integrated over all latitudes), there is a transition to anticorrelation around $r=0.9R_\odot$."
+" Through the bulk of the convection zone the high latitudes seem to be rotating faster during the solar activity maximum, and the equatorial regions rotate more slowly."," Through the bulk of the convection zone the high latitudes seem to be rotating faster during the solar activity maximum, and the equatorial regions rotate more slowly."
+" It is difficult to draw any definite conclusions about the temporal variation below the convection zone, however, because the errors are large."," It is difficult to draw any definite conclusions about the temporal variation below the convection zone, however, because the errors are large."
+ The amplitude of the temporal variation in the rotational kinetic energy integrated over the entire convection zone is of order of 3x1035 erg., The amplitude of the temporal variation in the rotational kinetic energy integrated over the entire convection zone is of order of $3\times10^{38}$ erg.
+ This implies a rate of variation whose, This implies a rate of variation whose
+planets.,planets.
+ These early successes have led to funding of a second generation of microlensing planet searches., These early successes have led to funding of a second generation of microlensing planet searches.
+ The first generation combines a wide-field search [or microlensing events with followup of the most promising events by small telescopes to search for planets., The first generation combines a wide-field search for microlensing events with followup of the most promising events by small telescopes to search for planets.
+ This vields 2-3 planets per vear., This yields 2-3 planets per year.
+ In the second stage. wide-fiekl cameras will continuously monitor about 16 deg?. which will vield dozens of planets per vear.," In the second stage, wide-field cameras will continuously monitor about 16 $\rm deg^2$, which will yield dozens of planets per year."
+ Note that while US astronomers plaved a major role in stage one. none (zero) of the roughly $440M required to build the second generation is coming [rom US sources.," Note that while US astronomers played a major role in stage one, none ) of the roughly 40M required to build the second generation is coming from US sources."
+ The third generation is a space-based xvide-field imager., The third generation is a space-based wide-field imager.
+ Detailed simulations show that it will improve both the mass limit and the number of detections by an order of magnitude relative to stage (wo., Detailed simulations show that it will improve both the mass limit and the number of detections by an order of magnitude relative to stage two.
+" Again. it is not my purpose here to review these simulations or designs. but to Gy to push the thinking of the committee ""out of the box”."," Again, it is not my purpose here to review these simulations or designs, but to try to push the thinking of the committee “out of the box”."
+ DE science will be best served by a mission (hat can be launched quickly aud so can obtain early results. stimulate new theoretical work. as well as new scientific ancl engineering ideas on how to proceed to the next step.," DE science will be best served by a mission that can be launched quickly and so can obtain early results, stimulate new theoretical work, as well as new scientific and engineering ideas on how to proceed to the next step."
+ No convincing argument can be made for attacking DE first by WL. ΟΔΟ. or SN.," No convincing argument can be made for attacking DE first by WL, BAO, or SN."
+ All have merits and demerits., All have merits and demerits.
+ llowever. a WL mission will simultaneously enable a search [or microlensing planets (hat will revolutionize the field of extra-solar planets.," However, a WL mission will simultaneously enable a search for microlensing planets that will revolutionize the field of extra-solar planets."
+Even though the present flow is turbulent. (he laminar approximation for the Stokes drag was used because (he grains were presumed to be smaller than the smallest turbulent eddies al all times.,"Even though the present flow is turbulent, the laminar approximation for the Stokes drag was used because the grains were presumed to be smaller than the smallest turbulent eddies at all times."
+ Therefore. unlike a macroscopic object in a turbulent flow. the erain would see an effectively laminar flow.," Therefore, unlike a macroscopic object in a turbulent flow, the grain would see an effectively laminar flow."
+ This approximation can be justified a posteriori: It turns out Chat only very close to the star. do the largest scale dust grains just before eravilational imstablitv become comparable to (he smallest scale turbulence.," This approximation can be justified a posteriori: It turns out that only very close to the star, do the largest scale dust grains just before gravitational instablity become comparable to the smallest scale turbulence."
+ We neglect the change in drag law that this would imply because (hose edcdies both are far less massive than (he grains and are traveling al velocities nearly equal to these dust erains at the onset. gravitational instability and so will have negligeable effects on the grain trajectories., We neglect the change in drag law that this would imply because those eddies both are far less massive than the grains and are traveling at velocities nearly equal to these dust grains at the onset gravitational instability and so will have negligeable effects on the grain trajectories.
+ The most uncertain aspect of our model lies in the sticking mechanism., The most uncertain aspect of our model lies in the sticking mechanism.
+ A value of p=1l. which represences efficiency. for collisions resulting in sticking. seems unlikely.," A value of $p=1$, which represences efficiency for collisions resulting in sticking, seems unlikely."
+ Considering p to be a gestalt parameter combining both all (he possible sticking and destruction results of collisions and examining Eq., Considering $p$ to be a gestalt parameter combining both all the possible sticking and destruction results of collisions and examining Eq.
+ 3 we can see that our time scale 1 depends inversely on p., \ref{main} we can see that our time scale $T$ depends inversely on $p$.
+ Comparison will sec., Comparison with sec.
+ 7.2 we see that decreasing e stronely decreases (he outer value for allowed by the heacdwincl induced infall. allhoueh this will only become important if p nears LO7.," 7.2 we see that decreasing $c$ strongly decreases the outer value for $R$ allowed by the headwind induced infall, although this will only become important if $p$ nears $10^{-2}$."
+ If we had a well developed sticking model we could caleulate the maxiniun collisional velocity before grain collisions result in destruction., If we had a well developed sticking model we could calculate the maximum collisional velocity before grain collisions result in destruction.
+ Accordingly. our ability to constrain a. either through disallowing grain destruction through high velocity collisions or through (ime scale arguments is limited by Che quality of the sticking model.," Accordingly, our ability to constrain $\alpha$, either through disallowing grain destruction through high velocity collisions or through time scale arguments is limited by the quality of the sticking model."
+ In our model we assume that all the dust grains share a common radius ry. and in section 2 we justify this assumption bv appealing to an eventual Gaussian specirum.," In our model we assume that all the dust grains share a common radius $r_d$, and in section 2 we justify this assumption by appealing to an eventual Gaussian spectrum."
+ As we saw in section 7.2. the differential rotation of the dust and σας discs caused by partial pressure support of the latter induces size dependent velocities in the former. (he radial one eiven bx Eq.," As we saw in section 7.2, the differential rotation of the dust and gas discs caused by partial pressure support of the latter induces size dependent velocities in the former, the radial one given by Eq."
+ 44. and an orbital one (Nakagawaetal. (1986))):, \ref{vr} and an orbital one \citet{Nakagawa86}) ):
+Ciureutlv. about 50 accretion-powered X-ray. pulsars are known (c.c.. Bildstenetal. 1997... VauParadijs 1995)).,"Currently, about 50 accretion-powered X-ray pulsars are known (e.g., \cite{batse}, , \cite{jvp95}) )."
+ The majority of these are located in binary svstenis (a few cases are thought to accrete from molecular clouds. e.9.. Corbet et al.," The majority of these are located in binary systems (a few cases are thought to accrete from molecular clouds, e.g., Corbet et al."
+ 1995)., 1995).
+ If optically identified. the companion is uxuallv a highauass star.," If optically identified, the companion is usually a high-mass star."
+ Iu these cases half are eiaut or πρσα! stars aud half are Oc or Be type stars., In these cases half are giant or supergiant stars and half are Oe or Be type stars.
+ Tuterestinely. he subset of those that are X-ray transient all have Oc or Be type companious.," Interestingly, the subset of those that are X-ray transient all have Oe or Be type companions."
+ This testifies in favor of the rotion that such stars loose matter iu an regular fashion (c.e.. vandeuIeuve&Rappaport 1987)).," This testifies in favor of the notion that such stars loose matter in an irregular fashion (e.g., \cite{Heuvel}) )."
+ From the vpical intrinsic bunimosities and distances of X-ray pulsars. it is expected that a relatively large population of faint (i... £ünter than ? 1) A-rav pulsars exists in our galaxy.," From the typical intrinsic luminosities and distances of X-ray pulsars, it is expected that a relatively large population of faint (i.e., fainter than $^{-9}$ ) X-ray pulsars exists in our galaxy."
+" lu fact. due to successful observations withGinga, the Rossi X-ray Timing Explorer aud this has been coufined (c.e.. Iiuugasaetal.1998... Isovainaetal.1990... Marsden: et al."," In fact, due to successful observations with, the Rossi X-ray Timing Explorer and this has been confirmed (e.g., \cite{AX J1820}, \cite{scutum}, Marsden et al."
+ 1998)., 1998).
+ Furthermore. I&oviuua et al. (," Furthermore, Koyama et al. ("
+1990) observed a concentration of pulsus iu the Seutum arm. very likely associated with a region of receut star formation.,"1990) observed a concentration of pulsars in the Scutum arm, very likely associated with a region of recent star formation."
+ We report the discovery of another faint transient X- pulsar using data accunuulated bv the Wide Field Cameras on board the BeppoSAX satellite., We report the discovery of another faint transient X-ray pulsar using data accumulated by the Wide Field Cameras on board the BeppoSAX satellite.
+ We preset the observations aud detection iu section 2.. the analysis of the pulsar signal in section 3.. aud the energy spectra in section [..," We present the observations and detection in section \ref{secobs}, , the analysis of the pulsar signal in section \ref{sec:period}, and the energy spectrum in section \ref{secspec}."
+ We discuss the implications of the details of the measurements for the nature of the source in section 5.., We discuss the implications of the details of the measurements for the nature of the source in section \ref{sec:discuss}.
+ The observations were performed with the Wide Field Camera (WFC) instrament (Jager et al., The observations were performed with the Wide Field Camera (WFC) instrument (Jager et al.
+ 1997)., 1997).
+ This instrument consists of two coded aperture cameras that point in opposite directions at the X-rav sky., This instrument consists of two coded aperture cameras that point in opposite directions at the X-ray sky.
+ They are located on the BeppoSAX N-zvav satellite (c.¢.. Boclla et al.," They are located on the BeppoSAX X-ray satellite (e.g., Boella et al."
+ 1997) which is in operation since April 1996., 1997) which is in operation since April 1996.
+ Each calucra is equipped with a multi2vire proportional xenon counter with a bandpass of 2 to 25 keV. The field of view is LO by LO square deerees of the skv)., Each camera is equipped with a multi-wire proportional xenon counter with a bandpass of 2 to 25 keV. The field of view is 40 by 40 square degrees of the sky).
+ The aneular resolution is5.. the accuracy to position a point source is 2 to 3 times better than that.," The angular resolution is, the accuracy to position a point source is 2 to 3 times better than that."
+ About of the observations with WEC are carried out in secondary mode., About of the observations with WFC are carried out in secondary mode.
+ In this mode. observations with the narrow-field instirmments on BeppoSAX dictate the oricutation of the satellite and. thus. the pointing direction of the WFC.," In this mode, observations with the narrow-field instruments on BeppoSAX dictate the orientation of the satellite and, thus, the pointing direction of the WFC."
+ The Cyenus feld is frequeutly covered by secoudary mode observations., The Cygnus field is frequently covered by secondary mode observations.
+ During the first 1.5 vr of the uuission the WECs were exposed to it during 90 days. with a fairly uniform time coverage.," During the first 1.5 yr of the mission the WFCs were exposed to it during 90 days, with a fairly uniform time coverage."
+ The total effective exposure time G.c.. when the Earth is outside the field of view) is ~2< 108s. Diving April 1997. a faint G.c.. for WFEC standards) transient was detected in secondary mode observations of the Cyveuus field. 7.1 degrees from) the bright X-rav source (νο A-3.," The total effective exposure time (i.e., when the Earth is outside the field of view) is $\sim2\times10^6$ s. During April 1997, a faint (i.e., for WFC standards) transient was detected in secondary mode observations of the Cygnus field, 7.4 degrees from the bright X-ray source Cyg X-3."
+" The error box of this transient is preseuted in fig. Ἐν,", The error box of this transient is presented in fig. \ref{figerrbox}.
+ A search of this boxwith the Simbad database gavea negative result on au identification with, A search of this boxwith the Simbad database gavea negative result on an identification with
+The external magnetic field of a neutron star is (relatively) easily inferred. from its spin-down rate. but its internal magnetic field is not directly observable.,"The external magnetic field of a neutron star is (relatively) easily inferred from its spin-down rate, but its internal magnetic field is not directly observable."
+ Phe main clue suggesting the existence of strong internal neutron star fields comes from the 1908 August 27 eiant Dare from the soft ganuna-ray repeater (SCAR) 1900|14 1999)..," The main clue suggesting the existence of strong internal neutron star fields comes from the 1998 August 27 giant flare from the soft gamma-ray repeater (SGR) 1900+14 \citep{fetal99,hetal99,metal99}."
+" The giant Hare. which released ~10°"" J of energy as N-rays. was accompanied. by a 2.3-fold. increase in the spin-down rate (Mazetsetal.1999:WoodsctThompson 2000)..."," The giant flare, which released $\sim\ 10^{37}$ J of energy as X-rays, was accompanied by a 2.3-fold increase in the spin-down rate \citep{metal99,wetal99,tetal00}. ."
+" To explain this. loka(2001) proposed that the Hare and the enhanced spin down were caused by a global reconfiguration of the internal magnetic field of ~LOM ""P. wellnove the external dipole field of 6.41027""EL."," To explain this, \citet{i01} proposed that the flare and the enhanced spin down were caused by a global reconfiguration of the internal magnetic field of $\sim 10^{13}$ T, wellabove the external dipole field of $6.4\times 10^{10}$."
+ Corsi&Owen(2011). generalised the loka(2001) caleulation (bv. allowing the toroidal fiel strength to change. as well as the moment of inertia) and concluded that an internal feld strength of ~107 a eave rise to the 1998 August 27 event. lower than the first estimate by loka(2001).. but still significantly higher than the οπονο external field.," \citet{co11} generalised the \citet{i01} calculation (by allowing the toroidal field strength to change, as well as the moment of inertia) and concluded that an internal field strength of $\sim 10^{12}$ T gave rise to the 1998 August 27 event, lower than the first estimate by \citet{i01}, but still significantly higher than the observed external field."
+ Stellar cllipicity can also be used to constrain the strength of a stars internal field 2011)..," Stellar ellipticity can also be used to constrain the strength of a star's internal field \citep{c02,detal09,aetal10,p11}."
+ Ht is well known that a strong magnetic field can deform a star (Chandrasekhar&Fermi 2011)..," It is well known that a strong magnetic field can deform a star \citep{cf53,f54,g72,k89,pm04,hetal08,metal11}. ."
+ Phe ellipticity 6 isroughly proxortional to the magnetic energv. (Cutler2002:Haskelletal.2008:Dall'Osso 2000).," The ellipticity $\epsilon$ isroughly proportional to the magnetic energy \citep{c02,hetal08,detal09}. ."
+. Neutron stars.," Neutron stars,"
+"that the sign of the coefficient K, can be either positive or negative.",that the sign of the coefficient $K_1$ can be either positive or negative.
+ But this difference has no effect on the conditions under which the disc and binary co- or counteralign., But this difference has no effect on the conditions under which the disc and binary co– or counteralign.
+" These are formally identical with the ones for the LT case derived by KLOP, with the binary angular momentum J, replacing the hole spin angular momentum J."," These are formally identical with the ones for the LT case derived by KLOP, with the binary angular momentum $\mathbi{J}_{\rm{b}}$ replacing the hole spin angular momentum $\mathbi{J}_{\rm{h}}$."
+ The process obviously has a different timescale specified by the different magnitude of Κ2., The process obviously has a different timescale specified by the different magnitude of $K_2$.
+" The same arguments as in KLOP now show that the magnitude J of the binary angular momentum remains constant, while the direction of J, aligns with the total angular momentum J»+ Ja, which is of course a constant vector."," The same arguments as in KLOP now show that the magnitude $J_{\rm{b}}$ of the binary angular momentum remains constant, while the direction of $\mathbi{J}_{\rm{b}}$ aligns with the total angular momentum $\mathbi{J}_{\rm{t}} = \mathbi{J}_{\rm{b}} + \mathbi{J}_{\rm{d}}$ , which is of course a constant vector."
+ During this process the magnitude of decreases because of dissipation (KLOP)., During this process the magnitude of $J_{\rm{d}}^2$ decreases because of dissipation (KLOP).
+ Counteralignment (0J}— π) occurs if and only if J;>J?.," Counteralignment $\theta \rightarrow \pi$ ) occurs if and only if $J_{\rm{b}}^2 >
+J_{\rm{t}}^2$."
+" By the cosine theorem so this is equivalent to Thus counteralignment of a binary and an external disc is possible, and requires So far in this paper we have avoided fully spelling out the meaning of the disc angular momentum Jy."," By the cosine theorem so this is equivalent to Thus counteralignment of a binary and an external disc is possible, and requires So far in this paper we have avoided fully spelling out the meaning of the disc angular momentum $\mathbi{J}_{\rm{d}}$."
+" This is complicated because the binary torque falls off very strongly with radius, and so a large contribution to the angular momentum in a distant part of the disc may be irrelevant to the alignment process, or affect this process in a time-dependent way (cf ?))."," This is complicated because the binary torque falls off very strongly with radius, and so a large contribution to the angular momentum in a distant part of the disc may be irrelevant to the alignment process, or affect this process in a time–dependent way (cf \citealt{LP2006}) )."
+ Sections 3 and 4 of KLOP discuss these questions in more detail., Sections 3 and 4 of KLOP discuss these questions in more detail.
+" Effectively Ja can be thought of as the disc angular momentum inside the warp radius, and therefore a time-dependent quantity."," Effectively $\mathbi{J}_{\rm d}$ can be thought of as the disc angular momentum inside the warp radius, and therefore a time–dependent quantity."
+" At early times J, is small, as only fraction of the total gas interacts with the binary."," At early times $J_{\rm{d}}$ is small, as only a fraction of the total gas interacts with the binary."
+" Counter-alignmenta may occur if 0>2/2, but at later times, as Jy grows and more gas is able to interact with the binary, alignment eventually happens (when Ja> 209)."," Counter–alignment may occur if $\theta > \pi/2$, but at later times, as $J_{\rm{d}}$ grows and more gas is able to interact with the binary, alignment eventually happens (when $J_{\rm{d}} > 2J_{\rm{b}}$ )."
+" So if 0>π/2, even for Jag>2Jy we expect ~2Js|cos6| of disc angular momentum to counteralign with the binary before the outer disc comes to dominate and enforce coalignment (cf. ?))."," So if $\theta
+>\pi/2$, even for $J_{\rm{d}} > 2J_{\rm{b}}$ we expect $\sim
+2J_{\rm{b}}|\cos\theta|$ of disc angular momentum to counteralign with the binary before the outer disc comes to dominate and enforce coalignment (cf. \citealt{LP2006}) )."
+" The typical timescale for co— or counter-alignment for a SMBH binary is where ἂν is the warp radius, Jy(R,) is the disc angular momentum within Αν, and v; is the vertical disc viscosity."," The typical timescale for co– or counter–alignment for a SMBH binary is where $R_{\rm{w}}$ is the warp radius, $J_{\rm{d}}(R_{\rm{w}})$ is the disc angular momentum within $R_{\rm{w}}$, and $\nu_2$ is the vertical disc viscosity."
+" This is identical to the formal expression for LT alignment of a spinning black hole if we replace the spin angular momentum J, with J, (cf ?)).", This is identical to the formal expression for LT alignment of a spinning black hole if we replace the spin angular momentum $J_{\rm{h}}$ with $J_{\rm{b}}$ (cf \citealt{SF1996}) ).
+" The warp radius is given by equating the precession time 1/Q,(R) to the vertical viscous time R?[v,.", The warp radius is given by equating the precession time $1/\Omega_{\rm{p}}(R)$ to the vertical viscous time $R^2/\nu_2$.
+ Inside this radius the precession timescale is short and the disc dissipates and co- or counter-aligns with the binary plane., Inside this radius the precession timescale is short and the disc dissipates and co– or counter–aligns with the binary plane.
+ Outside this radius the disc is not dominated by the precession and so maintains its misaligned plane., Outside this radius the disc is not dominated by the precession and so maintains its misaligned plane.
+ The connecting region therefore takes on a warped shape shown in 1., The connecting region therefore takes on a warped shape shown in .
+. As time passes the warp propagates outwards and co- or counter—aligns the entire disc with the binary plane., As time passes the warp propagates outwards and co– or counter–aligns the entire disc with the binary plane.
+" Approximating the disc angular momentum as with X the disc surface density and M=Μι+M5, and using the steady-state disc relation M=3zvX we find where we have also used Since γι«v; (7), a«Ry and M»«Mi, we see that alignment takes place on a timescale shorter than the mass growth of the central accretor(s)."," Approximating the disc angular momentum as with $\Sigma$ the disc surface density and $M = M_1 + M_2$, and using the steady–state disc relation $\dot M = 3\pi\nu\Sigma$ we find where we have also used Since $\nu_1 < \nu_2$ \citep{PP1983}, $a \ll R_{\rm{w}}$ and $M_2 < M_1$, we see that alignment takes place on a timescale shorter than the mass growth of the central accretor(s)."
+" The timescale (15)) is directly analogous to the expression for alignment under the LT precession, where a,« is the Kerr spin parameter and R, the Schwarzschild radius of1 the spinning hole."," The timescale \ref{talign2}) ) is directly analogous to the expression for alignment under the LT precession, where $a_* < 1$ is the Kerr spin parameter and $R_{\rm{s}}$ the Schwarzschild radius of the spinning hole."
+" Evaluating A, in the two cases we find Thus in general, provided we assume that the ratio v;/v» is similar in the two cases and that the hole spin is not rather small (a.«(R,/a)(Mo M), then the binary-disc alignment is rather faster than the corresponding process for spinning black holes."," Evaluating $R_{\rm{w}}$ in the two cases we find Thus in general, provided we assume that the ratio $\nu_1/\nu_2$ is similar in the two cases and that the hole spin is not rather small $a_* <
+(R_{\rm{s}}/a)^{1/2}(M_2/M_1)$ ), then the binary–disc alignment is rather faster than the corresponding process for spinning black holes."
+ Our result has significant consequences for SMBH binaries., Our result has significant consequences for SMBH binaries.
+" For random orientations, shows that initial disc angles leading to alignment occur significantly more frequently than those giving counteralignment only if J4>29. ("," For random orientations, shows that initial disc angles leading to alignment occur significantly more frequently than those giving counteralignment only if $J_{\rm d} > 2J_{\rm b}$. ("
+"In the LT case this fact leads to a slow spindown of the hole, because retrograde accretion has a larger effect on the spin, ?..)","In the LT case this fact leads to a slow spindown of the hole, because retrograde accretion has a larger effect on the spin, \citealt{Kingetal2008}. .)"
+ A number of studies (?;; ?;; 25; ?)) have shown that prograde external discs are rather inefficient in shrinking SMBH binaries and solving the last parsec problem., A number of studies \citealt{AN2005}; \citealt{MM2008}; \citealt{Cuadraetal2009}; \citealt{Lodatoetal2009}) ) have shown that prograde external discs are rather inefficient in shrinking SMBH binaries and solving the last parsec problem.
+ This is essentially because of resonances within the disc., This is essentially because of resonances within the disc.
+" In contrast, the slightly rarer retrograde events have a much stronger effect onthe binary."," In contrast, the slightly rarer retrograde events have a much stronger effect onthe binary."
+" These rapidly produce a counterrotating but coplanar accretion disc external to the binary, which has no resonances."," These rapidly produce a counterrotating but coplanar accretion disc external to the binary, which has no resonances."
+ We note that ?) , We note that \citet{Nixonetal2011} 
+because {μον were interested mainiv in the dvnamical evolution of the PWN.,because they were interested mainly in the dynamical evolution of the PWN.
+ In this paper. we revisit a spectral evolution model of PWNe. paving attention to the issues mentioned above.," In this paper, we revisit a spectral evolution model of PWNe, paying attention to the issues mentioned above."
+ We do not take into account effects of the spatial structure of the PWN because it is somewhat costly and too detailed to discuss the whole radiation spectrum., We do not take into account effects of the spatial structure of the PWN because it is somewhat costly and too detailed to discuss the whole radiation spectrum.
+ In our model. the PWN is simply treated as an expanding uniform sphere.," In our model, the PWN is simply treated as an expanding uniform sphere."
+ The enerev inside the PWN is injected from the pulsar spin-down energy. which is divided between the magnetic field aud (he relativistic particles with a constant ratio.," The energy inside the PWN is injected from the pulsar spin-down energy, which is divided between the magnetic field and the relativistic particles with a constant ratio."
+ Our simple model can describe the observed basic features., Our simple model can describe the observed basic features.
+ Of course. many details are not treated in (his simplified model. such as the filamentary structures and the spatial variations of photon indices.," Of course, many details are not treated in this simplified model, such as the filamentary structures and the spatial variations of photon indices."
+ We study the spectral evolution of the Crab Nebula as the first application of our model., We study the spectral evolution of the Crab Nebula as the first application of our model.
+ The Crab Nebula can be used as a calibrator of our model when it will be applied to other PWNe in future., The Crab Nebula can be used as a calibrator of our model when it will be applied to other PWNe in future.
+ In section ??.. we describe our model of the PWN evolution.," In section \ref{model}, we describe our model of the PWN evolution."
+ In section ??.. we apply (his model to the Crab Nebula.," In section \ref{crab}, we apply this model to the Crab Nebula."
+" In section οοι, discussions and conclusions are mace."," In section \ref{discussion_conclusion}, discussions and conclusions are made."
+ For the calculation of the spectral evolution. we need to specilv the evolution of the magnetic field and the particle distribution funcGon.," For the calculation of the spectral evolution, we need to specify the evolution of the magnetic field and the particle distribution function."
+ Here. we describe the assumptions of our moclel.," Here, we describe the assumptions of our model."
+as lieh temperature isophotes within the vvoluue.,as high temperature isophotes within the volume.
+" While the large temperature thresholds reduce the confusion along the line of sight. these necessarily bias the resultant cloud catalogs to the warmest. deusest regions within the molecular iuterstellar πουπια,"," While the large temperature thresholds reduce the confusion along the line of sight, these necessarily bias the resultant cloud catalogs to the warmest, densest regions within the molecular interstellar medium."
+ aand lnaeing observations of targeted clouds demonstrate that most of the molecular mass resides within the extended. low column deusitv lines of sight (Carpenter. Sucll Schloerb 1995: ever. Carpenter Ladd 1996).," and imaging observations of targeted clouds demonstrate that most of the molecular mass resides within the extended, low column density lines of sight (Carpenter, Snell Schloerb 1995; Heyer, Carpenter Ladd 1996)."
+ Such low colin deusityv regious iu the iuner Calaxy are simply not accessible for the analysis of cloud. properties., Such low column density regions in the inner Galaxy are simply not accessible for the analysis of cloud properties.
+ Iu coutrast. the outer Galaxy provides a less confusing view of the molecular iuterstellar iiediuuu.," In contrast, the outer Galaxy provides a less confusing view of the molecular interstellar medium."
+ Bevoucd the solar circle. there is no blending of cussion from widely separated clouds aloug the line of sight.," Beyond the solar circle, there is no blending of emission from widely separated clouds along the line of sight."
+ Therefore. molecular reeious can be identified at lower eas column deusities from which more represeutative elobal properties can be derived.," Therefore, molecular regions can be identified at lower gas column densities from which more representative global properties can be derived."
+ This property las been exploited in a series of investigations by Braud Wouterloot (1991. 1995).," This property has been exploited in a series of investigations by Brand Wouterloot (1994, 1995)."
+ While these studies provide a scusitive. high resolution perspective of individual clouds aud the distribution of molecular regions in the far outer Galaxy. the results are necessarily biased toward clouds associated with star formation.," While these studies provide a sensitive, high resolution perspective of individual clouds and the distribution of molecular regions in the far outer Galaxy, the results are necessarily biased toward clouds associated with star formation."
+ The FCRAO CO Survey of the outer Calaxy. provides an opportunity to study the equilibrium| state of iiolecular clouds wider varving conditions (Hever 1998)., The FCRAO CO Survey of the outer Galaxy provides an opportunity to study the equilibrium state of molecular clouds under varying conditions (Heyer 1998).
+" The survey searched forο J=1-0 emission within. a330o deg? 2pfield siupled every wwith a FWIIM beam size of15"".", The survey searched for$^{12}$ CO J=1-0 emission within a330 $^2$ field sampled every with a FWHM beam size of.
+. The range is -153 to LO ssuaupled every 0.81 wwith a resolution of 0.98ο, The range is -153 to 40 sampled every 0.81 with a resolution of 0.98.
+"ν, The median main beam sensitivity (19) per chaunel is 0.9 Ix. Iu this contribution. we present results from a decomposition of the outer Galaxy Survey into discrete objects."," The median main beam sensitivity $\sigma$ ) per channel is 0.9 K. In this contribution, we present results from a decomposition of the outer Galaxy Survey into discrete objects."
+ Thuninosity. size. aud line width distributions are determined from the euseiible of ideutified objects located in the Perseus ari aud far outer Galaxy.," luminosity, size, and line width distributions are determined from the ensemble of identified objects located in the Perseus arm and far outer Galaxy."
+ Iu 63. we reexanuue the Larson scaling relationships with the cloud catalog extracted from the Survey aud with chuups from a sinülar decomposition of aand oobservations of several targeted giant molecular clouds.," In $\S$ 3, we reexamine the Larson scaling relationships with the cloud catalog extracted from the Survey and with clumps from a similar decomposition of and observations of several targeted giant molecular clouds."
+ To isolate discrete regions of CO emission from the large data cube. we lave adopted the definition of a nolecular cloud used. by. previous investigations (Solomon 1987: Scoville 1987: Sodroski 19011.," To isolate discrete regions of CO emission from the large data cube, we have adopted the definition of a molecular cloud used by previous investigations (Solomon 1987; Scoville 1987; Sodroski 1991)."
+ That is. a discrete molecular region is identified as a closed topological surface within the?obVra data cube at a eiven threshold of auteuna temperature.," That is, a discrete molecular region is identified as a closed topological surface within the $l-b-V_{LSR}$ data cube at a given threshold of antenna temperature."
+ In this study. the limiting threshold is 1.118 πα beam cluperature scale) or 1.56 where o is the median ruis value of autenua temperatures in the Survey (Wever 1998).," In this study, the limiting threshold is 1.4 K (main beam temperature scale) or $\sigma$ where $\sigma$ is the median rms value of antenna temperatures in the Survey (Heyer 1998)."
+ The threshold value is suticicutly low to provide a more complete accounting of the flux within he data cube as compared to the inner Galaxy surveys. while large enough to exclude misideutificatious of nolecular regions due to statistical noise.," The threshold value is sufficiently low to provide a more complete accounting of the flux within the data cube as compared to the inner Galaxy surveys, while large enough to exclude misidentifications of molecular regions due to statistical noise."
+ Detailed descriptions of the cloud decomposition aud the calculation of cloud properties are provided im Appendix A. The decomposition of the FCRAO CO Survey of the Outer Calaxy at a limiting threshold of Tyjp-11.1 IK. vields 10156 objects., Detailed descriptions of the cloud decomposition and the calculation of cloud properties are provided in Appendix A. The decomposition of the FCRAO CO Survey of the Outer Galaxy at a limiting threshold of 1.4 K yields 10156 objects.
+"Each object is described by position ceutroids (ωνδρ 0.). velocity width.óc... a kinematic distance. D. assuming purely circular motious and a flat rotation curveον, Galactoceutrie radius. R4. Z Leight. CO luninosity.Leo. aud a peak autenua temperature within the suface.T,.","Each object is described by position centroids $l_c,b_c,v_c$ ), velocity width, a kinematic distance, D, assuming purely circular motions and a flat rotation curve, Galactocentric radius, $_{gal}$, z height, CO luminosity, and a peak antenna temperature within the surface,."
+. The geometry of au object is described by major aud minor axis diameters. hijas. hj. and a position angle. 0. of the," The geometry of an object is described by major and minor axis diameters, , , and a position angle, $\theta$ , of the"
+We expect Chis result to be reasonably accurate only when for (he reasons mentioned above.,We expect this result to be reasonably accurate only when $0.9 \le T/T_c \le 1$ for the reasons mentioned above.
+" In that case the formula can be expanded just below 7, as follows: tel — In2- ) - 4) (5", In that case the formula can be expanded just below $T_c$ as follows: 1 - 2 - ) - 4 ].
+9) Note that the effect. ofn the(225) confinement. scale is to reduce the shear relaxation time compared to the conformal limit., Note that the effect of the confinement scale is to reduce the shear relaxation time compared to the conformal limit.
+ Numerically. (he value of T;. can be estimated as follows.," Numerically, the value of $T_c$ can be estimated as follows."
+ The mass of the paneson in the hard wall model is lound by solving a particular wave equation whose solutions are Bessel functions., The mass of the $\rho$ -meson in the hard wall model is found by solving a particular wave equation whose solutions are Bessel functions.
+ A boundary condition requires that οπρμμ)=0 from which one clecluces that T;2innf1-95022109 MeV. This is just a characteristic temperature. not the critical temperature of a phase (transition. since one may define it somewhat differently [13]..," A boundary condition requires that $J_0(m_{\rho}L^2/r_{\rm min}) = 0$ from which one deduces that $T_c \approx m_{\rho}/7.556 \approx 102$ MeV. This is just a characteristic temperature, not the critical temperature of a phase transition, since one may define it somewhat differently \cite{softwallD}."
+ The soft wall model [19] was developed to improve upon the hard wall model., The soft wall model \cite{softwall} was developed to improve upon the hard wall model.
+ In particular. it leads to linear Regee trajectories wherein the radial excitations of (he vector and axial-vector meson spectra have mass-scquared being linear in the radial quantum number ». in substantial agreement wilh data.," In particular, it leads to linear Regge trajectories wherein the radial excitations of the vector and axial-vector meson spectra have mass-squared being linear in the radial quantum number $n$, in substantial agreement with data."
+ It can be obtained by adding a dilaton field to the usual AdS metric., It can be obtained by adding a dilaton field to the usual AdS metric.
+" The piece of the action that is relevant lor the meson mass spectra is 7 7 +om Επ) 45meson where o is the dilaton field and F7"" and Fe” are the field strength tensors for the left and right handed gauge fields."," The piece of the action that is relevant for the meson mass spectra is = d^5 x (F_L^2 + F_R^2) where $\phi$ is the dilaton field and $F_L^{\mu \nu}$ and $F_R^{\mu
+\nu}$ are the field strength tensors for the left and right handed gauge fields."
+ In our coordinates. the dilaton prolile which leads to the Regge behavior is ο)=cL'/i?. where e is a constant which can be determined by fitting the meson spectrum.," In our coordinates, the dilaton profile which leads to the Regge behavior is $\phi(r) = c L^4/r^2$, where $c$ is a constant which can be determined by fitting the meson spectrum."
+ If we compute the quantities 77/5 and 7. using the soft wall model of |19].. the results are the sanie as Eqs. (35))," If we compute the quantities $\eta/s$ and $\tau$ , using the soft wall model of \cite{softwall}, the results are the same as Eqs. \ref{ConformalEta}) )"
+ and (5)) because the metric is exactly κάδο., and \ref{ConformalTau}) ) because the metric is exactly $AdS_5$.
+ There is an alternative Formulation of the soft wall., There is an alternative formulation of the soft wall.
+ Instead of adding a nontrivialdilaton. one keeps (he dilaton constant while deforming themetric away [rom Ads.," Instead of adding a nontrivialdilaton, one keeps the dilaton constant while deforming themetric away from $AdS_5$ ."
+Landau damping: as à result svnchrotron emission could be insignificant.,Landau damping; as a result synchrotron emission could be insignificant.
+ Medvedev Loch (1999) found more reassuring results., Medvedev Loeb (1999) found more reassuring results.
+ Relativistic wo-streanr magnetic instability can. they claim. generate stable randomly oriented strong magnetic field in the plane of the collisionless shock front. so svnchrotron emission. is possible.," Relativistic two-stream magnetic instability can, they claim, generate stable randomly oriented strong magnetic field in the plane of the collisionless shock front, so synchrotron emission is possible."
+ These contradictory results show our ignorance of the possible acceleration. process. and of how (or if) a stable magnetic field is created there.," These contradictory results show our ignorance of the possible acceleration process, and of how (or if) a stable magnetic field is created there."
+ In the standard mocel his ignorance is “parameterized  through ¢ and eg: the simplistic assumption is made that a persistent. magnetic ield makes these parameters uniform in space and time., In the standard model this ignorance is ”parameterized “ through $\epsilon_e$ and $\epsilon_B$: the simplistic assumption is made that a persistent magnetic field makes these parameters uniform in space and time.
+" An observation that maybe suggests a time dependence of c, aud cg is the independent determination of these parameters in GRB 970508 afterglow by Wijers ancl Galama (1999) at 12 davs and Frail. Waxman and Kulkarni (2000) at ~I vear: while the former found ο.=0.12 and cg=0.089 the latter inferred e,£cg0.0."," An observation that maybe suggests a time dependence of $\epsilon_e$ and $\epsilon_B$ is the independent determination of these parameters in GRB 970508 afterglow by Wijers and Galama (1999) at 12 days and Frail, Waxman and Kulkarni (2000) at $\sim 1$ year: while the former found $\epsilon_e=0.12$ and $\epsilon_B=0.089$ the latter inferred $\epsilon_e\simeq \epsilon_B=0.5$."
+ These uncertainties. coupled with the apparent problems with the standard model. motivate us to explore the consequences. for afterglow emission. of a. dilferent scenario. where the propagation of the magnetic field to large scale is cdisfavored.," These uncertainties, coupled with the apparent problems with the standard model, motivate us to explore the consequences for afterglow emission of a different scenario, where the propagation of the magnetic field to large scale is disfavored."
+ In this contest we particularly focus on the possibility that external density could be larger. keeping the total οσον to a reasonable value. without loosing a eencral agreement with data.," In this contest we particularly focus on the possibility that external density could be larger, keeping the total energy to a reasonable value, without loosing a general agreement with data."
+ In fact. at this stage. we deal with sparse set of measurements. almost never simultaneous at all wavelengths.," In fact, at this stage, we deal with sparse set of measurements, almost never simultaneous at all wavelengths."
+ Vhese data can be modeled in dilferent frameworks. (e.g. GRB 970508 (Frail et al 2000. Chevalier Li 2000). GRB 980519 (Frail et al 2000b. Chevalier Li 19999. GRB 000301C (Bereer et al 2000. Li Chevalier 2001). GRB 991208 (Calama et al 2000. Li Chevalier 2001 and GRB 000418 (Bereer οἱ al 20012).," These data can be modeled in different frameworks, (e.g. GRB 970508 (Frail et al 2000, Chevalier Li 2000), GRB 980519 (Frail et al 2000b, Chevalier Li 1999), GRB 000301C (Berger et al 2000, Li Chevalier 2001), GRB 991208 (Galama et al 2000, Li Chevalier 2001 and GRB 000418 (Berger et al 2001))."
+ ‘This analysis is meant to investigate how robust are the current estimation of shock and density. parameters., This analysis is meant to investigate how robust are the current estimation of shock and density parameters.
+ In this paper we present this more general theory for the GRB afterglow emission (82): our conclusion (83) is that a qualitative comparison of this model with cata suggests the possibility that a shorter magnetic lengthscale can be involved. alfecting the estimates of parameters. but only a quantitative modeling of clata can give the conclusive ÜLSWOLD.," In this paper we present this more general theory for the GRB afterglow emission 2); our conclusion 3) is that a qualitative comparison of this model with data suggests the possibility that a shorter magnetic lengthscale can be involved, affecting the estimates of parameters, but only a quantitative modeling of data can give the conclusive answer."
+ The assumption of a post shock magnetic field. persisting during all the expansion time has three main consequences for the svnchrotron (S) emission and spectrum., The assumption of a post shock magnetic field persisting during all the expansion time has three main consequences for the synchrotron (S) emission and spectrum.
+ First the emitting region linear dimension is A2c5foy. where Foy is the expansion time in the lab frame: this enters in the calculation. of the observed. peak Εαν £5.," First the emitting region linear dimension is $R\simeq c \times t_{exp}$, where $t_{exp}$ is the expansion time in the lab frame; this enters in the calculation of the observed peak flux $F_p$."
+ Then the observed. cooling frequency κ. ds computed. using the Lorenz factor. σοι of the electrons that cool raciativelv on a timescale equal to the remnant age.," Then the observed cooling frequency $\nu_c$ is computed using the Lorenz factor, $\gamma_c$, of the electrons that cool radiatively on a timescale equal to the remnant age."
+ Finally the Iengthscale for synchrotron self-absorption has to be compared with the dimension of the fireball in order to compute the observed absorption frequency αι, Finally the lengthscale for synchrotron self-absorption has to be compared with the dimension of the fireball in order to compute the observed absorption frequency $\nu_a$.
+" “Phe only. break frequeney which is not allected by any assumption about the extension. of the magnetic field. is the observed. peak frequeney η. corresponding to the minimum random Lorentz factor 5: it depends crucially only on the fractions of internal energy eiven to electrons. c,. and given to the magnetic field ον."," The only break frequency which is not affected by any assumption about the extension of the magnetic field, is the observed peak frequency $\nu_m$, corresponding to the minimum random Lorentz factor $\gamma_m$: it depends crucially only on the fractions of internal energy given to electrons, $\epsilon_e$, and given to the magnetic field $\epsilon_b$."
+ Phroughout the entire emission volume electrons. cool also via inverse Compton (1€ on the svnchrotron photons (Self Synchrotron Compton: SSC)., Throughout the entire emission volume electrons cool also via inverse Compton (IC) on the synchrotron photons (Self Synchrotron Compton: SSC).
+ This cooling process dominates the electrons cooling if the Compton parameter )—LiefLs (4€ and S luminosity ratio) is greater than one., This cooling process dominates the electrons cooling if the Compton parameter $Y=L_{IC}/L_S$ (IC and S luminosity ratio) is greater than one.
+" This happens if c,« and current estimates find in most objects enfe,~LO1.107. (Mijers CGalama 1999. Ciranot et al 1999. POL)."," This happens if $\epsilon_b \ll \epsilon_e$ and current estimates find in most objects $\epsilon_b/\epsilon_e\sim 10^{-1}-10^{-2}$, (Wijers Galama 1999, Granot et al 1999, PK01)."
+ If the magnetic field behind the shock persists for an average timescale of £j<£s. its lengthscale is A2es£e Ro ," If the magnetic field behind the shock persists for an average timescale of $t_b<t_{exp}$, its lengthscale is $\lambda\simeq c \times
+t_b<R$ ."
+ln principle A could. depend on time. however we clo not have anv consolidated: understanding of this process. therefore for simplicity we will assume throughout this paper that ΑΙ=9 ds constant during. the afterglow phase.," In principle $\lambda$ could depend on time, however we do not have any consolidated understanding of this process, therefore for simplicity we will assume throughout this paper that $\lambda/R=\delta$ is constant during the afterglow phase."
+ At cach time the emitting region is then divided in two zones (see Fig 11)., At each time the emitting region is then divided in two zones (see Fig \ref{fig:car}) ).
+ Phe first zone (region A) is the laver with dimension A just behind the shock front. where there are newly accelerated. electrons. and. magnetic. Ποιά: here electrons emit. via S and SSC: the second zone (region D) of linear dimension A. is filled by electrons that have just cooled by S and SSC emissions and that now Compton scatter the radiation coming from the first laver.," The first zone (region A) is the layer with dimension $\lambda $ just behind the shock front, where there are newly accelerated electrons and magnetic field: here electrons emit via S and SSC; the second zone (region B) of linear dimension $R-\lambda$, is filled by electrons that have just cooled by S and SSC emissions and that now Compton scatter the radiation coming from the first layer."
+ We suppose that the newly accelerated electrons are advected with the How: they produce svnchrotron radiation only for a short time (with a cooling. break therefore only at very high energies) in region A near the shock front anc then they emit most of the LC radiation while clownstreaming through region D. On the other hand i£ the electrons cillused Lrecly. so that they could move back into the magnetized region after having lost most of their energy via IC. then the svnchrotron spectrum would have a Ilower-energy break.," We suppose that the newly accelerated electrons are advected with the flow; they produce synchrotron radiation only for a short time (with a cooling break therefore only at very high energies) in region A near the shock front and then they emit most of the IC radiation while downstreaming through region B. On the other hand if the electrons diffused freely, so that they could move back into the magnetized region after having lost most of their energy via IC, then the synchrotron spectrum would have a lower-energy break."
+ This is probably less Likely therefore we concentrate on the first. case., This is probably less likely therefore we concentrate on the first case.
+ The total observed spectrum then results from the sum of region A and region D spectra., The total observed spectrum then results from the sum of region A and region B spectra.
+would have 6 star clusters beyoud LOO pc.,would have 6 star clusters beyond 400 pc.
+ Model N.77.11.5.2200 has the same cutoll radius and therefore the same values of 3 as model 1100., Model 200 has the same cutoff radius and therefore the same values of $\beta$ as model 100.
+ Due to the broader initial distribution of star cluster in. M_77_11.5_2200 the enclosed mass aud the effective radius are slightly lower thau in inodel 1100., Due to the broader initial distribution of star cluster in 200 the enclosed mass and the effective radius are slightly lower than in model 100.
+ The results demoustrate that next to the CC mass and the CC Plumiuer radius also the cutoll racius aud the exact distribution of star clusters in a CC are key parameters for exteuded moclels. cousiderably inllueucing the structural parameters of the merger objects.," The results demonstrate that next to the CC mass and the CC Plummer radius also the cutoff radius and the exact distribution of star clusters in a CC are key parameters for extended models, considerably influencing the structural parameters of the merger objects."
+" The enclosed inasses were converted into absolute W-magnituces to allow for direct comparison with the observed data. usiug the formula where. Myi, = L853 mag is the absolute solar V. iuagnitude. Mj the euclosed mass of the merger object aud ic the inass-to-liglit ratio."," The enclosed masses were converted into absolute $V$ -magnitudes to allow for direct comparison with the observed data, using the formula where, $M_{\rm V,solar}$ = 4.83 mag is the absolute solar $V$ magnitude, $M_{\rm encl}$ the enclosed mass of the merger object and $\frac{M}{L_V}$ the mass-to-light ratio."
+ We use a mass-to-light ratio of 2.0540.50. as determined by Batunearcdtοἱal.(2009) for 22119.," We use a mass-to-light ratio of $\pm$ 0.50, as determined by \cite{baumgardt} for 2419."
+ Apparent V-maenituces can by caleulatec by adding the distance modulus of 22119 (uM) = 19.60 (Ripepietal.2007) and a V-bane extinction of Ay = 0.25 using a reddeuiug of E(B-V) = 0.08 (Ripepietal., Apparent V-magnitudes can by calculated by adding the distance modulus of 2419 $_0$ = 19.60 \citep{ripepi} and a V-band extinction of $A_{\rm V}$ = 0.25 using a reddening of E(B-V) = 0.08 \citep{ripepi}.
+2YT).. Figure 11. shows au exemplary coutour plot of the merger object of moclel 1100 projected onto the sky using Galactic coordinates., Figure \ref{contourplot} shows an exemplary contour plot of the merger object of model 100 projected onto the sky using Galactic coordinates.
+ The lowest contour line corresponds to 32 mag 7., The lowest contour line corresponds to 32 mag $^{-2}$.
+ At this low surface brightuess. 22119 is detectable up to radii of about ccorrespoudiug to 115 pe.," At this low surface brightness, 2419 is detectable up to radii of about corresponding to 415 pc."
+ Observed photometry. however. covers only a region within the inner ool 22119. while surface densities need to be derived [rom star counts iu the outer parts.," Observed photometry, however, covers only a region within the inner of 2419, while surface densities need to be derived from star counts in the outer parts."
+ Ripepietal.(2007) found evidence for stars associated with 22119 up to radii of15%., \cite{ripepi} found evidence for stars associated with 2419 up to radii of.
+ Figure 12 shows the effective radius. reg. versus the absolute ruagnitude. My. ofthe observations aud the models as given in Table 1. and 3.. respectively.," Figure \ref{obssims} shows the effective radius, $r_{\rm eff}$ , versus the absolute magnitude, $M_{\rm V}$, of the observations and the models as given in Table \ref{tbl-1} and \ref{tbl-2}, respectively."
+ Due to the scatter of the observed parameters of 22119 and the turnover in the ellective radii of the models as described iu Section ??.. à largeuumber of mocels are compatible with the observed parameters of 22119.," Due to the scatter of the observed parameters of 2419 and the turnover in the effective radii of the models as described in Section \ref{massandsize}, a largenumber of models are compatible with the observed parameters of 2419."
+ Only extended models with initial CC masses of ALS° = 1.0 x 10° and compact models with Aloo = 2.0 x 10° are clearly incompatible with observations., Only extended models with initial CC masses of $M^{\rm CC}$ = 1.0 $\times$ $^{6}$ and compact models with $M^{\rm CC}$ = 2.0 $\times$ $^{6}$ are clearly incompatible with observations.
+ A more detailed comparison between models aud observations can be achieved. using surface brightness profiles., A more detailed comparison between models and observations can be achieved using surface brightness profiles.
+ Bellazziui(2007) compiled a surface brightuess profile out to projected radii of 177755., \cite{bellazzini} compiled a surface brightness profile out to projected radii of 5.
+" However. as he used radii between aand 900"".. which are well inside 22119. as a referencefield to estimate the background level."," However, as he used radii between and , which are well inside 2419, as a referencefield to estimate the background level,"
+"Recent theoretical studies based on cosmological sinulations sueecst that the first stars are formed in dark matter (DAL) halos with mass 10?)10°AT, at τν30 (Teemarketal.L997:Yoshidactal. 2003).","Recent theoretical studies based on cosmological simulations suggest that the first stars are formed in dark matter (DM) halos with mass $10^5-10^6 \ \mathrm{M_{\odot}}$ at $z \sim 20-30$ \citep{tegmark97, yoshida03}."
+. Tn such a cosmological “unimihalo”. the eas cools and condeuses by molecular hydrogen cooling to fori) a star-forming eas cloud.," In such a cosmological “minihalo”, the gas cools and condenses by molecular hydrogen cooling to form a star-forming gas cloud."
+ The eravitationallv uustable cloud further coutracts. and finally a proto-star is born inside it (Ommukai&Nishi1998:Yoshidaetal. 2008).," The gravitationally unstable cloud further contracts, and finally a proto-star is born inside it \citep{omukai98, yoshida08}."
+. A unique characteristic of ornuordiali star formation in the standard cold dark matter model is that the star and its parent eas cloud are embedded at the center o| the hos dark halo., A unique characteristic of primordial star formation in the standard cold dark matter model is that the star and its parent gas cloud are embedded at the center of the host dark halo.
+ Thus the formation and evolution of primordia stars nv be mich affected x DAL dynamically., Thus the formation and evolution of primordial stars may be much affected by DM dynamically.
+" The nature of DM remains still ποσα, but the best candidates are thought to be Woeaklv Tuteracting Massive Particles. WIMPs."," The nature of DM remains still unknown, but the best candidates are thought to be Weakly Interacting Massive Particles, WIMPs."
+ WIMP DAL such as neutralinos must have a laree sclaunililation cross-section iu order to be a dominant conrponeut of DAL in the present-day universe., WIMP DM such as neutralinos must have a large self-annihilation cross-section in order to be a dominant component of DM in the present-day universe.
+ DM inunihilatiou produces an enormous amount of energy. essentially equal to the rest mass energy of the annilulated DAL particles.," DM annihilation produces an enormous amount of energy, essentially equal to the rest mass energy of the annihilated DM particles."
+ If DAL deusity can be very large in a primordial gas cloud aud in a primordial star. the annihilation energv would act as an efficient heat input.," If DM density can be very large in a primordial gas cloud and in a primordial star, the annihilation energy would act as an efficient heat input."
+ DM. aunililatiou could even supply μποτ eneregv to support selt-evavitation of a star., DM annihilation could even supply sufficient energy to support self-gravitation of a star.
+ This is indeed a recently proposed new stellar phase. in which stars are powered bv DAL annihilation euergv imustead of that of the nuclear fusion (Spolvarctal.2008).," This is indeed a recently proposed new stellar phase, in which stars are powered by DM annihilation energy instead of that of the nuclear fusion \citep{spolyar08}."
+. Iereafter we call such a star star., Hereafter we call such a star ”.
+ Spolvarctal.(2008) sugecst that. in a collapsing primordial eas cloud. DAL aunililation heating can win overthe radiative cooling. effectively halts further collapse of the cloud.," \citet{spolyar08} suggest that, in a collapsing primordial gas cloud, DM annihilation heating can win over the radiative cooling, effectively halts further collapse of the cloud."
+ Natarajanetal.(2009)— study the evolution of the DM density profile during the first star formation usiug cosmological simulations., \citet{natarajan09} study the evolution of the DM density profile during the first star formation using cosmological simulations.
+ They show that the DM density profile indeed becomes very steep. as ich as expected by adiabatie contraction models.," They show that the DM density profile indeed becomes very steep, as much as expected by adiabatic contraction models."
+ There have been many previous studies o- dark stars., There have been many previous studies on dark stars.
+ While most of thoi assiunue coustaut stellar mass models (loccoetal.2008:Taosoetal.2008:Yoonctal," While most of them assume constant stellar mass models \citep{iocco08, taoso08, yoon08}."
+ 2008).. Spolvaretal.(2009). aud ποιαetal.(2009). study dark star evolution with eas mass accretion., \citet{spolyar09} and \citet{umeda09} study dark star evolution with gas mass accretion.
+ By using an approximated stellar structure model. follow dark star evolutions with adiabatically coutracted DM (AC'DAL).," By using an approximated stellar structure model, follow dark star evolutions with adiabatically contracted DM (ACDM)."
+ They show that the final stellar mass becomes as large as 1000NE... Uined, They show that the final stellar mass becomes as large as $1000 \ \mathrm{M_{\odot}}$.
+aetal.(2009). perform detailed. dynamically self-cousisteut stellar evolution calculations up to exavitatioual core-collapse.," \citet{umeda09} perform detailed, dynamically self-consistent stellar evolution calculations up to gravitational core-collapse."
+ They show that the final uass can be as huge as 105ALL or more. if the star captures DAL efficiently.," They show that the final mass can be as large as $10^{4-5} \ \mathrm{M_{\odot}}$ or more, if the star captures DM efficiently."
+ However. the DAL capture rate itself is uncertain because it depends on the unkuown DM-baryou scattering cross section and also ou the ambient DAL density diving the dark star evolution.," However, the DM capture rate itself is uncertain because it depends on the unknown DM-baryon scattering cross section and also on the ambient DM density during the dark star evolution."
+ Thus the proposed ‘orlnation of iuterauediate massive black holes ypears ο occur onlv iu very particular cases., Thus the proposed formation of inter-mediate massive black holes appears to occur only in very particular cases.
+ An attractive feature of the dark star models is hat the possibility for the birth of massive stars. and black holes (BIIS) as the eud of such stars. in he carly universe.," An attractive feature of the dark star models is that the possibility for the birth of massive stars, and black holes (BHs) as the end of such stars, in the early universe."
+ Super-assive DIIs (SAIBUs) existed. already at iον6. when the ageo of the universe is less than one billion vears old. but he formation mechanisia of such early SMDIIS relmaius unknown.," Super-massive BHs (SMBHs) existed already at $z \sim 6$, when the age of the universe is less than one billion years old, but the formation mechanism of such early SMBHs remains unknown."
+ Gravitational collapse of a very nassive dark star nüght provide a solution to this xoblemi., Gravitational collapse of a very massive dark star might provide a solution to this problem.
+ We explore a possibility that the first stars formed in DM halos can become very niassive x being powered by DM annihilation., We explore a possibility that the first stars formed in DM halos can become very massive by being powered by DM annihilation.
+ To this eud. it is portant to determine the final state of the dark star pliase exactly aud to deteruiue the final stellar mass.," To this end, it is important to determine the final state of the dark star phase exactly, and to determine the final stellar mass."
+ Iun this paper. we study the evolution of an accretiue dark star using a detailed stellar evolution code as in Uinedaetal.(2009). but with the ACDAL annihilation.," In this paper, we study the evolution of an accreting dark star using a detailed stellar evolution code as in \citet{umeda09}
+ but with the ACDM annihilation."
+ We calculate the density profile of ACDAL using an analytical method aud determine the DM annihilation rate., We calculate the density profile of ACDM using an analytical method and determine the DM annihilation rate.
+ We hen follow the evolution of the dark star which erows im mass by gas accretion., We then follow the evolution of the dark star which grows in mass by gas accretion.
+ The dark star model has two parameters effectivelv: the eas Lnuass accretion rate and DAL particle mass., The dark star model has two parameters effectively; the gas mass accretion rate and DM particle mass.
+ We calculate several models with different set of parameters aud investigate the effects on the final stellar mass., We calculate several models with different set of parameters and investigate the effects on the final stellar mass.
+ Especidbve we clarity the effect of the differeut mass accretion rates. which were not explicitly shown in a similar work by Spolvaral. (2009).," Especially, we clarify the effect of the different mass accretion rates, which were not explicitly shown in a similar work by \citet{spolyar09}."
+. The res of the paper is organized as follows., The rest of the paper is organized as follows.
+ In section 2. we introduce our muuerical calculation and explain how we implement DAT annihilation iu the stellar evolution.," In section 2, we introduce our numerical calculation and explain how we implement DM annihilation in the stellar evolution."
+ Sec., Sec.
+ 3 shows the results of our dark star models aud we diseuss the implications, 3 shows the results of our dark star models and we discuss the implications
+The predictions for cosmological parameter errors presented in this paper use the Fisher matrix formalism.,The predictions for cosmological parameter errors presented in this paper use the Fisher matrix formalism.
+ The Fisher matrix gives us the lower bound on the accuracy with which we can estimate model parameters from a given cata set (Fisher1935:Tegmark.Tavlor.&Leavens1997:Witching&Amara 2009).," The Fisher matrix gives us the lower bound on the accuracy with which we can estimate model parameters from a given data set \citep*{Fisher1935, Tegmark:1997, Kitching:2009}."
+. In. calculating forecast survey errors. we are implicitly making assumptions about theparameter set (scethediscussionofnestedmodelsinHeavens.Witching.&Verde 2007).," In calculating forecast survey errors, we are implicitly making assumptions about theparameter set \citep*[see the discussion of nested models in][]{Heavens:2007}."
+. We want to know whether our constraints are robust against variations in the parameterisation of the cosmological model., We want to know whether our constraints are robust against variations in the parameterisation of the cosmological model.
+ This is of particular importance when clark energy constraints are considered. because of the degeneracies with other parameters (see Fig. 1)).," This is of particular importance when dark energy constraints are considered, because of the degeneracies with other parameters (see Fig. \ref{pk_deg}) )."
+ The forecast parameter precision is improved. bv combining independent experiments., The forecast parameter precision is improved by combining independent experiments.
+ Using this technique. joint Constraints or error forecasts can be obtained by combining weak lensing with other observational techniques.," Using this technique, joint constraints or error forecasts can be obtained by combining weak lensing with other observational techniques."
+ The Fisher matrix for the shear power spectrum is given by (Hu&Jain2004):: = where the sum is over bands of multipole { of width Af. ‘Tr is the trace. ancl fa is the fraction of sky covered by the survey.," The Fisher matrix for the shear power spectrum is given by \citep{Hu:2004}: = where the sum is over bands of multipole $\ell$ of width $\Delta \ell$, $\mr{Tr}$ is the trace, and $f_\mr{sky}$ is the fraction of sky covered by the survey."
+ Equation ?? assumes the likelihod obeys a Gaussian distribution with zero mean., Equation \ref{eq_10} assumes the likelihod obeys a Gaussian distribution with zero mean.
+" The observed. power spectra for each pair 7. of redshift bins are written as the sum of the lensing and noisej spectra: The derivative matrices are given by where p, is the vector of parameters in the theoretical niocel."," The observed power spectra for each pair $i,j$ of redshift bins are written as the sum of the lensing and noise spectra: The derivative matrices are given by where $p_\alpha$ is the vector of parameters in the theoretical model."
+ In order to quantify the potential for à survey to constrain dark energv parameters. we use the Figure of Merit. as defined by the Dark Energy Task Force (Albrechtetal. 2006):: In this article. together with lensing-onlv constraints. we also include joint lensing and CAIB constraints.," In order to quantify the potential for a survey to constrain dark energy parameters, we use the Figure of Merit, as defined by the Dark Energy Task Force \citep{DETF}: In this article, together with lensing-only constraints, we also include joint lensing and CMB constraints."
+ To combine constraints fron cifferent probes we add: the respective Fisher matrices: Fyvin=ΈιμιΕσυ., To combine constraints from different probes we add the respective Fisher matrices: $\bf{F}_\mr{joint}=\bf{F}_\mr{lensing}+\bf{F}_\mr{CMB}$.
+ For our CMD priors. we use the forthcoming. mission as our survey.," For our CMB priors, we use the forthcoming mission as our survey."
+ Phe Fisher matrix is calculated ollowing Rassatetal.(2008).. which estimates errors using information from the temperature ancl I2. moce »olarisation (i.e. TTE. EI. and TI. with the help of the xibliclv available code (Lewis.Challinor.&Lasenbvy 2000).," The Fisher matrix is calculated following \citet{Rassat:2008}, which estimates errors using information from the temperature and E mode polarisation (i.e. TT, EE, and TE), with the help of the publicly available code \citep*{Lewis:2000}."
+. We conservatively do not use information from B modes. and only use the 143 Gllz channel. assuming other requencies will be used for the foreground removal.," We conservatively do not use information from B modes, and only use the 143 GHz channel, assuming other frequencies will be used for the foreground removal."
+ This is conservative compared. to other priors in. the iterature., This is conservative compared to other priors in the literature.
+ More details are given in Rassatetal.(2008.Ap-»encdix B).., More details are given in \citet[][Appendix B]{Rassat:2008}.
+" Phe full parameter set for the calculation is: (Opp.wo.wa.OQ,On.mo.Newboos.nusa.7}."," The full parameter set for the calculation is: $\{\Omega_\mr{DE},\,w_0,\,w_a,\,\Omega_m,\,\Omega_b,\,m_\nu,\, N_\nu,\,h,\,\sigma_8,\,n_s,\,\alpha, \,\tau\}$."
+ We use the same central values as for our weak lensing calculations. as described above. with a fiducial value for the reionisation optical depth 7=0.09 and subsequently marginalize over 7.," We use the same central values as for our weak lensing calculations, as described above, with a fiducial value for the reionisation optical depth $\tau=0.09$ and subsequently marginalize over $\tau$."
+" We consider neutrino parameters as O57 and IN, and use a Jacobian to translate this into constraints on my, and UN.", We consider neutrino parameters as $\Omega_\nu h^2$ and $N_\nu$ and use a Jacobian to translate this into constraints on $m_\nu$ and $N_\nu$.
+ Weak lensing measurements are adfected bv systematic clleets which οσο the precision on cosmological parameters and introduce bias (see ltéfréeier.2003. and Schneider2006. for a review)., Weak lensing measurements are affected by systematic effects which reduce the precision on cosmological parameters and introduce bias (see \citealt{Refregier:2003a} and \citealt{Schneider:2006} for a review).
+ Lo this section we discuss some of these ellects., In this section we discuss some of these effects.
+ Intrinsic correlations can contaminate the lensing signal., Intrinsic correlations can contaminate the lensing signal.
+ Solutions include using tomography or 3D. lensing. which decouple the long-clistance line-of-sight. effects. [rom the the physical proximity of the galaxies (seee.g.Wing&Schneider2003:Bricle&Wine2007:Witchingetal. 2008b).," Solutions include using tomography or 3D lensing, which decouple the long-distance line-of-sight effects from the the physical proximity of the galaxies \citep[see e.g.][]{King:2003, Bridle:2007, Kit2008b}."
+. Using this approach. a nulling technique for shear-intrinsic nu has been proposed by Joachimi&Schneider(2008.TWO).," Using this approach, a nulling technique for shear-intrinsic ellipticity has been proposed by \citet{Joachimi:2008, Joachimi:2009}."
+ Measurement systematics are due to PSE effects. ]xaiseretal. 1995)., Measurement systematics are due to PSF effects \citep[see][]{KSB1995}.
+. The results presented. in. Dridleοἱal.(2009) indicate that the required accuracy will be reached for the next generation of all-sky weak lensing surveys., The results presented in \citet{Bridle:2009} indicate that the required accuracy will be reached for the next generation of all-sky weak lensing surveys.
+ Redshift) cistribution svsteniatics. which lead to an uncertaintv in the median galaxy redshift. cause an uncertainty in the amplitude of the matter power spectrum (Llu&Tegmark1999)..," Redshift distribution systematics, which lead to an uncertainty in the median galaxy redshift, cause an uncertainty in the amplitude of the matter power spectrum \citep{Hu:1999b}."
+ The problem of photometric redshift systematics can be met given a number of galaxies in the spectroscopic calibration seuuple of 103107 (Maetal.2006:Amara&ltélrégier 2007).," The problem of photometric redshift systematics can be met given a number of galaxies in the spectroscopic calibration sample of $10^4-10^5$ \citep{Ma:2006, AR2007}."
+. Finally there are theoretical uncertainties on the shape of the matter power spectrum., Finally there are theoretical uncertainties on the shape of the matter power spectrum.
+ The existing non-linear corrections to the matter power spectrum (Peacock&Doclds1996:Ma1998). are only accurate to about 1054 and disagree with one another to this level in the non-linear régeime Lluterer 2001).," The existing non-linear corrections to the matter power spectrum \citep{Peacock:1996ys, Ma:1998} are only accurate to about $10\%$ and disagree with one another to this level in the non-linear réggime \citep[see][]{Huterer:2001}."
+. Newer prescriptions such as those bv Smithetal.(2003). olfer more accurate predictions particularly for non-ACDAL cosmological models., Newer prescriptions such as those by \citet{Smithetal} offer more accurate predictions particularly for $\Lambda\mr{CDM}$ cosmological models.
+" Phe error in the non-linear part may still be significant if elfect of massive neutrinos is included (seee.g.Saito.Takada.&""Tarüva2008).", The error in the non-linear part may still be significant if effect of massive neutrinos is included \citep*[see e.g.][]{Saito:2008}.
+. Current semi-analvtical models need to be improved to match the degree of statistical accuracy expected for future weak lensing surveys., Current semi-analytical models need to be improved to match the degree of statistical accuracy expected for future weak lensing surveys.
+ The solution is to run a suite of N-body rav-tracing simulations (seee.g.White&Vale2004:Llutereretal. 2009). ," The solution is to run a suite of $N$ -body ray-tracing simulations \citep[see e.g.][]{White:2004, Huterer:2005, Hilbert:2009, Sato:2009, Teyssier:2009}. ."
+In our Fisher matrix formalism. the error forecast on cach parameter depends on the sensitivity of the weak lensing observation to changes in the matter power spectrum.," In our Fisher matrix formalism, the error forecast on each parameter depends on the sensitivity of the weak lensing observation to changes in the matter power spectrum."
+ In order to probe the elfect of the dilferent. parameters on the matter power spectrum.we consider the fractional change in the non-linear matter power spectrum (A). defined as the change in P(A) with respect to the ficlucial 2A) ia. when one parameter at a time is varied from its fiducial value:," In order to probe the effect of the different parameters on the matter power spectrum,we consider the fractional change in the non-linear matter power spectrum $P(k)$ , defined as the change in $P(k)$ with respect to the fiducial $P(k)_\mr{fid}$ , when one parameter at a time is varied from its fiducial value:"
+which there is not enough information to determine if a signal is real.,which there is not enough information to determine if a signal is real.
+ Primarily this excludes periods longer than the time—span of the data. but early and late in the time series it is also affected by the padding. hence the cone-like shape.," Primarily this excludes periods longer than the time–span of the data, but early and late in the time series it is also affected by the padding, hence the cone–like shape."
+More specifically. the wavele power is expected to drop as we approach the edge of the data. due to the infinite support of the Morlet wavelet function.,"More specifically, the wavelet power is expected to drop as we approach the edge of the data, due to the infinite support of the Morlet wavelet function."
+ Fig., Fig.
+ 4aa depicts the wavelet power plot of the optical ligh curve in Fig., \ref{fig-power}a a depicts the wavelet power plot of the optical light curve in Fig.
+ | following re-sampling to a one day interval and interpolation of missing data as described above., \ref{fig-lc} following re-sampling to a one day interval and interpolation of missing data as described above.
+ The vertica dotted lines represent the beginning of an outburst as determined from our parametrisation approach., The vertical dotted lines represent the beginning of an outburst as determined from our parametrisation approach.
+ All the peaks in the wavele power spectrum are associated with an outburst., All the peaks in the wavelet power spectrum are associated with an outburst.
+ A similar feature can be located for the year 1945. where Oppenheimer&Mattei(1993). suggested an additional outburst.," A similar feature can be located for the year 1945, where \citet{Oppenheimer1993} suggested an additional outburst."
+ However. the wavele power is not as strong as all the other outbursts due to the lack of data during the brightest part of the outburst (as we found in Section 3. if an outburst had indeed occurred. it started during the seasonal gap).," However, the wavelet power is not as strong as all the other outbursts due to the lack of data during the brightest part of the outburst (as we found in Section \ref{sec-decline} if an outburst had indeed occurred, it started during the seasonal gap)."
+ The most prominent peak in the wavelet power spectrum of Fig., The most prominent peak in the wavelet power spectrum of Fig.
+ daa is linked with the 1967 outburst., \ref{fig-power}a a is linked with the 1967 outburst.
+ This has to do with the fact that during the last stage of the magnitude decline of the outburst the magnitude becomes fainter than 12. which might be an indication that the 1967 outburst was more powerful than the others observed.," This has to do with the fact that during the last stage of the magnitude decline of the outburst the magnitude becomes fainter than 12, which might be an indication that the 1967 outburst was more powerful than the others observed."
+ Note that the features we observe are not symmetric about the beginning of each outburst., Note that the features we observe are not symmetric about the beginning of each outburst.
+ This asymmetry is because the wavelet power is influenced by the data before and after the outbursts., This asymmetry is because the wavelet power is influenced by the data before and after the outbursts.
+ The signal for the 2006 outburst seems to be weaker than the rest of the outbursts because part of it is outside of the COI., The signal for the 2006 outburst seems to be weaker than the rest of the outbursts because part of it is outside of the COI.
+ This means that the wavelet power at larger scales will decrease as we approach the edges because we pad the end of the time series with the mean value., This means that the wavelet power at larger scales will decrease as we approach the edges because we pad the end of the time series with the mean value.
+ A natural response of the wavelet power to discontinuities in the data (such as those seen at outburst in Fig., A natural response of the wavelet power to discontinuities in the data (such as those seen at outburst in Fig.
+ dau) is a peak in the power at the time of the discontinuity., \ref{fig-power}a a) is a peak in the power at the time of the discontinuity.
+ We decided to remove the first 143 days of each outburst and undertake the same analysis., We decided to remove the first 143 days of each outburst and undertake the same analysis.
+ The removed data have been replaced with the mean of the time series. as presented in Fig.," The removed data have been replaced with the mean of the time series, as presented in Fig."
+ 4bb. We tested the robustness of the results for different choices for replacement data. including a sinusoid of 50 day period and the residuals of subtracting the model determined in Section 3..," \ref{fig-power2}b b. We tested the robustness of the results for different choices for replacement data, including a sinusoid of 50 day period and the residuals of subtracting the model determined in Section \ref{sec-decline}."
+ Introducing different artiticial values for the removed data did not affect the outcome of this analysis other than to insignificantly change the detailed distribution of power in the wavelet power spectrum., Introducing different artificial values for the removed data did not affect the outcome of this analysis other than to insignificantly change the detailed distribution of power in the wavelet power spectrum.
+ Fig., Fig.
+ 4bb shows the resultant wavelet power spectrum., \ref{fig-power2}b b shows the resultant wavelet power spectrum.
+ There are peaks in the wavelet power spectrum at a period range between 600 and SOO days that are associated with each of the 1945. 1985 and 2006 pre-outburst phases.," There are peaks in the wavelet power spectrum at a period range between 600 and 800 days that are associated with each of the 1945, 1985 and 2006 pre-outburst phases."
+ This implies that these three outbursts are most similar regarding the light curve., This implies that these three outbursts are most similar regarding the light curve.
+ The fact that part of the 2006 peak of the wavelet power is outside of the COT does not suggest that the wavelet power is completely wrong., The fact that part of the 2006 peak of the wavelet power is outside of the COI does not suggest that the wavelet power is completely wrong.
+ Since we pad with the mean value of the time series at the edges. we expect the wavelet power near the edges at large scales to be less than it otherwise would be.," Since we pad with the mean value of the time series at the edges, we expect the wavelet power near the edges at large scales to be less than it otherwise would be."
+ This means that if we had more years of data. then the 2006 wavelet power peak would be entirely inside the COT and appear with enhanced wavelet power.," This means that if we had more years of data, then the 2006 wavelet power peak would be entirely inside the COI and appear with enhanced wavelet power."
+ The 1967 outburst is associated with a feature in the period range 500 and 1100 days. whereas the 1958 outburst does not give any signiticant peak associated with the pre-outburst.," The 1967 outburst is associated with a feature in the period range 500 and 1100 days, whereas the 1958 outburst does not give any significant peak associated with the pre–outburst."
+ Thus we believe that the peak for 1958 in the wavelet power spectrum in Fig., Thus we believe that the peak for 1958 in the wavelet power spectrum in Fig.
+ Ibb. when the outburst data is retained.is due only to the data discontinuity.," \ref{fig-lc}b b, when the outburst data is retained,is due only to the data discontinuity."
+ It is clear that all of the other outbursts (apart from the 1958) are linked with a specific form of wavelet power variation roughly around a period of 700—800 days which appears before the outburst begins. and so cannot be due to the discontinuity.," It is clear that all of the other outbursts (apart from the 1958) are linked with a specific form of wavelet power variation roughly around a period of 700–800 days which appears before the outburst begins, and so cannot be due to the discontinuity."
+ Last but not least. the pre- phase and the post-outburst phase (i.e. after 143 days from the outburst) of the light curve seem to follow the same pattern in all cases.," Last but not least, the pre-outburst phase and the post-outburst phase (i.e. after 143 days from the outburst) of the light curve seem to follow the same pattern in all cases."
+ This is another indication why we should remove the outbursts from the analysis. as the peak for the wavelet power for the 1958 outburst in Fig.daa is created only by the data discontinuity.," This is another indication why we should remove the outbursts from the analysis, as the peak for the wavelet power for the 1958 outburst in \ref{fig-power}a a is created only by the data discontinuity."
+ Therefore. the 1958 and 1967 outbursts might have progressed differently and may provide a key to the outburst mechanisms.," Therefore, the 1958 and 1967 outbursts might have progressed differently and may provide a key to the outburst mechanisms."
+ In addition the 1958 outburst was only 9 years before the 1967 outburst. and also the 1958 outburst was not associated with any fluctuation around that range.," In addition the 1958 outburst was only 9 years before the 1967 outburst, and also the 1958 outburst was not associated with any fluctuation around that range."
+towards the west. dominate at high frequencies.,"towards the west, dominate at high frequencies."
+ This implies that and OH trace lines-of-sight that are different from those traced by CH3OH: and OH absorb against sources E and F and methanol absorption should mainly arise toward sources A. B. and E. The similarity of the spectral lineshapes requires the presence of extended molecular complexes that are at least as extended as the distance between sources B and F. ~50 mmas or ppc.," This implies that and OH trace lines-of-sight that are different from those traced by $_3$ OH; and OH absorb against sources E and F and methanol absorption should mainly arise toward sources A, B, and E. The similarity of the spectral lineshapes requires the presence of extended molecular complexes that are at least as extended as the distance between sources B and F, $\sim$ mas or pc."
+ At GGHz. continuum sources À. B. E. and F contribute only mmJy or to to the total flux density (Middelberg et al.," At GHz, continuum sources A, B, E, and F contribute only mJy or to to the total flux density (Middelberg et al."
+ 2007)., 2007).
+ If all of the methanol absorption is viewed against these components. the optical depths of the two absorption features are about 3-4 times larger than the optical depths given in Table 2..," If all of the methanol absorption is viewed against these components, the optical depths of the two absorption features are about 3–4 times larger than the optical depths given in Table \ref{table2}."
+" After having observed the J, = δι — 69 A transition of A-type methanol and having estimated line parameters and optical depths. physical and chemical implications still have to be discussed."," After having observed the $J_{\rm k}$ = $_1$ $\rightarrow$ $_0$ $^+$ transition of A-type methanol and having estimated line parameters and optical depths, physical and chemical implications still have to be discussed."
+ The lowest levels of the «20 ladder of A-type methanol have lower energies than those in the neighboring k-ladders that are connected by allowed radiative transitions., The lowest levels of the $k$ =0 ladder of A-type methanol have lower energies than those in the neighboring k-ladders that are connected by allowed radiative transitions.
+" As à consequence. the energy levels of the £=0 ladder tend to be overpopulated and. in the absence of a strong radiation field. one expects enhanced absorption (also called ""over-cooling"" or ""anti-inverston) in the /, = δι — 69 A transition."," As a consequence, the energy levels of the $k$ =0 ladder tend to be overpopulated and, in the absence of a strong radiation field, one expects enhanced absorption (also called “over-cooling” or “anti-inversion”) in the $J_{\rm k}$ = $_1$ $\rightarrow$ $_0$ $^+$ transition."
+ This is verified by the statistical equilibrium caleulations of Walmsley et al. (, This is verified by the statistical equilibrium calculations of Walmsley et al. (
+1988) and Leurini et al. (,1988) and Leurini et al. (
+2004) and describes a Class I methanol maser environment.,2004) and describes a Class I methanol maser environment.
+ The situation is dramatically different for sources in the vicinity of an intense FIR field resulting from warm dust (7 > KK) heated by an embedded high-mass protostellar object surrounded by cooler gas., The situation is dramatically different for sources in the vicinity of an intense FIR field resulting from warm dust $T_{\rm d}$ $>$ K) heated by an embedded high-mass protostellar object surrounded by cooler gas.
+ Here. intense pumping. predominantly to the first torsionally excited state via radiation around 304m. and subsequent decay determine the level populations and lead to strong GGHz maser emission (e.g.. Cragg et al.," Here, intense pumping, predominantly to the first torsionally excited state via radiation around $\mu$ m, and subsequent decay determine the level populations and lead to strong GHz maser emission (e.g., Cragg et al."
+ 2005)., 2005).
+ The mere fact that we observed absorption and not emission already tells us that pumping processes leading to widespread maser emission are not the dominant excitation mechanism for methanol toward the central region of 33079., The mere fact that we observed absorption and not emission already tells us that pumping processes leading to widespread maser emission are not the dominant excitation mechanism for methanol toward the central region of 3079.
+ While this does not exclude the presence of GGHz masers (a single JJy maser at D=Skkpe would show a flux density of only mmJy at the distance to 33079). absorption dominates. presumably because it is more widespread.," While this does not exclude the presence of GHz masers (a single Jy maser at $D$ kpc would show a flux density of only mJy at the distance to 3079), absorption dominates, presumably because it is more widespread."
+ This is remarkable because the nuclear region of 33079 appears to be more active than that of our own Galactic centre region. where deep absorption is also observed (Menten 1991).," This is remarkable because the nuclear region of 3079 appears to be more active than that of our own Galactic centre region, where deep absorption is also observed (Menten 1991)."
+ With the measured peak apparent optical depths (Table 2)). with a source covering factor f. ~ 0.25—-0.30 (end of Sect.5.2). and in the absence of strong excitation by FIR photons. we can estimate the methanol column density along the line-of- toward 33079.," With the measured peak apparent optical depths (Table \ref{table2}) ), with a source covering factor $f_{\rm c}$ $\sim$ 0.25–0.30 (end of Sect.5.2), and in the absence of strong excitation by FIR photons, we can estimate the methanol column density along the line-of-sight toward 3079."
+ There are. however. two important unknown parameters. the kinetic temperature and the density of the gas.," There are, however, two important unknown parameters, the kinetic temperature and the density of the gas."
+ With these values in the range of Των = KK and (Hz) = 100-109 cem? and the Large Velocity Gradient code kindly supplied by S. Leurint (see Leurini et al., With these values in the range of $T_{\rm kin}$ = K and $n$ $_2$ ) = $^3$ $^6$ $^{-3}$ and the Large Velocity Gradient code kindly supplied by S. Leurini (see Leurini et al.
+" 2004). we find that A-type methanol column densities. N(A-CH:OH). between 3x10' and 6x10"" cem""? produce absorption at the observed levels."," 2004), we find that A-type methanol column densities, $N$ $_3$ OH), between $\times$ $^{13}$ and $\times$ $^{15}$ $^{-2}$ produce absorption at the observed levels."
+ The main uncertainty Hes with the density. with high densities yielding low column densities and vice versa.," The main uncertainty lies with the density, with high densities yielding low column densities and vice versa."
+ Our calculations generally find excitation temperatures <I KK. Thus the 5; — 6) Av line indeed is over-cooled., Our calculations generally find excitation temperatures $\ll$ K. Thus the $_1$ $\rightarrow$ $_0$ $^+$ line indeed is over-cooled.
+ We have detected the GGHz transition of methanol. one of the most prolific Galactic maser lines. for the first time in à source located well beyond the Magellanic Clouds.," We have detected the GHz transition of methanol, one of the most prolific Galactic maser lines, for the first time in a source located well beyond the Magellanic Clouds."
+ The spectrum obtained toward the starburst and H2O megamaser galaxy 33079 shows two absorption components with line-to-continuum ratios of about 1/50 and 1/150., The spectrum obtained toward the starburst and $_2$ O megamaser galaxy 3079 shows two absorption components with line-to-continuum ratios of about 1/50 and 1/150.
+ This is the signature of a Class I environment. where the absence of a strong infrared radiation field inhibits the inversion of the level populations.," This is the signature of a Class I environment, where the absence of a strong infrared radiation field inhibits the inversion of the level populations."
+" Instead. the line is characterised by inversion"" or ""overcooling""."," Instead, the line is characterised by ``anti-inversion'' or “overcooling”."
+" Aost of the absorption likely occurs toward the nuclear continuum sources A. B. and E. With no information on the density and kinetic temperature of the gas. the A-methanol column density is poorly defined and may range between a few times 10"" and a few times 10° cem""."," Most of the absorption likely occurs toward the nuclear continuum sources A, B, and E. With no information on the density and kinetic temperature of the gas, the A-methanol column density is poorly defined and may range between a few times $^{13}$ and a few times $^{15}$ $^{-2}$."
+"4 Interferometric measurements will be needed to reliably convert ""apparent"" into “true” optical depths.", Interferometric measurements will be needed to reliably convert “apparent” into “true” optical depths.
+ The presence of GGHz CH;OH absorption hot only toward the central region of the Galaxy but also toward the more active spiral galaxy 33079 suggests that such gas can be found in other extragalactic sources as well., The presence of GHz $_3$ OH absorption not only toward the central region of the Galaxy but also toward the more active spiral galaxy 3079 suggests that such gas can be found in other extragalactic sources as well.
+ To reveal the physical parameters of the interstellar medium traced by the GGHz line. the detection of additional methanol lines would be desirable.," To reveal the physical parameters of the interstellar medium traced by the GHz line, the detection of additional methanol lines would be desirable."
+" Suitable candidates are the 25—3., E GGHz). «|-δη E GGHz). and 7;—σι A GGHz) transitions that should also be mapped. for comparison. toward the Galactic centre region."," Suitable candidates are the $_0-3_{-1}$ E GHz), $_{-1}-3_0$ E GHz), and $_0-6_1$ $^+$ GHz) transitions that should also be mapped, for comparison, toward the Galactic centre region."
+»»pulations are separate are deceptive and should not be aken to indicate that these objects are discrete families of eravitationallv bound stellar systems.,populations are separate are deceptive and should not be taken to indicate that these objects are discrete families of gravitationally bound stellar systems.
+ Although this lack of evident distinction is a set-back or attempts to differcutiate the populations aud probe »»teutial differcuces in their formation and evolution. it is an advantage to those efforts to build a simple set of models for galaxies.," Although this lack of evident distinction is a set-back for attempts to differentiate the populations and probe potential differences in their formation and evolution, it is an advantage to those efforts to build a simple set of models for galaxies."
+ Star clusters can provide an anchor for models of galaxies. cipirically definimus he stellar 1nass-to-light ratios or. alternatively. the zero xut of the FAL," Star clusters can provide an anchor for models of galaxies, empirically defining the stellar mass-to-light ratios or, alternatively, the zero point of the FM."
+ The latter would enable us to use he FM to determine the mass-to-light ratio within +... o a level of precision comparable to that obtained with detailed dvnamical modeling or strong eravitational ensine modeling. for any galaxy with known V.τον aud I..," The latter would enable us to use the FM to determine the mass-to-light ratio within $r_e$, to a level of precision comparable to that obtained with detailed dynamical modeling or strong gravitational lensing modeling, for any galaxy with known $V, r_e$ , and $I_e$."
+ We retun fo the question of the defining characteristics of a galaxy., We return to the question of the defining characteristics of a galaxy.
+ Among the stellar systems that we study; we have not identified a criterion that we consider a discriminating feature.," Among the stellar systems that we study, we have not identified a criterion that we consider a discriminating feature."
+ Despite the popularity of the dark matter halo criterion. there exist “galaxics” that may not have a dark matter halo (tidaldurgalax-ies:Barnes&Ioeruquist1992) or. at least. a dominant dark matter halo Qiltra-compactdwarfgalaxies:Mieskeetal.," Despite the popularity of the dark matter halo criterion, there exist “galaxies"" that may not have a dark matter halo \citep[tidal dwarf galaxies;][]{barnes} or, at least, a dominant dark matter halo \citep[ultra-compact dwarf galaxies;][]{mieske}."
+ 2008).. In. fact. from state-of-the-art. optical data alone. it is difücult to cemonstrate that even eiaut ellipticals contain dark matter (Cappellaotal.2006).," In fact, from state-of-the-art optical data alone, it is difficult to demonstrate that even giant ellipticals contain dark matter \citep{cappellari06}."
+". An alternative. aud more easily observable. classification is the relative couceutration of huuimous barvous to dark iatter,"," An alternative, and more easily observable, classification is the relative concentration of luminous baryons to dark matter."
+ Iu such a scheme. systems compose! solely of barvous are not distinct from those with dark halos aud whose observed kinematics are dominated bv barvous.," In such a scheme, systems composed solely of baryons are not distinct from those with dark halos and whose observed kinematics are dominated by baryons."
+ While this scheme makes it harder to connect observations with dark matter simulations. the reality is that. for svsteius whose barvous are extremely concentrated. we cannot determine whether a halo exists or what its.. properties mught be.," While this scheme makes it harder to connect observations with dark matter simulations, the reality is that, for systems whose baryons are extremely concentrated, we cannot determine whether a halo exists or what its properties might be."
+ Our primary cipirical result is that the existing description of the Fundamental Manifold (Zaritskyetal.WwKS) ds valid over au even larger parameter range than previously demonstrated and that now immcludes stellar clusters., Our primary empirical result is that the existing description of the Fundamental Manifold \citep{zzg} is valid over an even larger parameter range than previously demonstrated and that now includes stellar clusters.
+ The existence of a EM for the ecutive famuly of stellar svstemis implies 1) that the objects are virialized and 2) that there is lauited variation in the overall structure., The existence of a FM for the entire family of stellar systems implies 1) that the objects are virialized and 2) that there is limited variation in the overall structure.
+ While the former is wusurprisine. the latter iuplicitlv places strong Πες on how the svstenis form and how the Iuninous baryous settle.," While the former is unsurprising, the latter implicitly places strong limits on how the systems form and how the luminous baryons settle."
+ By extending the validity of the FM. or any scaling relation. over a larger range of paraiueter space and a wider set of objects; we are demonstrating that the physics that coustrains the resulting barvou aud dark matter distributions in stellar systenis Is more eenceral than previously thought.," By extending the validity of the FM, or any scaling relation, over a larger range of parameter space and a wider set of objects, we are demonstrating that the physics that constrains the resulting baryon and dark matter distributions in stellar systems is more general than previously thought."
+" In other words, whereas one might not have been surprised by the nuplication of the Fundamental Plane that all cllipticals have related structural properties. it is certainly a puzzle to fud that spirals. dwarf spheroidals and now stellar clusters do too."," In other words, whereas one might not have been surprised by the implication of the Fundamental Plane that all ellipticals have related structural properties, it is certainly a puzzle to find that spirals, dwarf spheroidals, and now stellar clusters do too."
+ While the structure of one stellar pAsystem. ids nof a shuple scaling of another. there existsyo a transform function that relates the properties of stellar svstenis.," While the structure of one stellar system is not a simple scaling of another, there exists a transform function that relates the properties of stellar systems."
+" Specifically. we have preseuted a fitting fuuction for Y, that is valid across all stellar systems."," Specifically, we have presented a fitting function for $\Upsilon_e$ that is valid across all stellar systems."
+" The ecuerality of this transform function inplies that 1) the structure of one type of stellar system is uot a distinct result of a process particular to that set of svstenis. such as violent relaxation for elliptical galaxies. but is mstead the result of an interplay of all processes responsible for the generic settling of barvous in eravitational poteutial wells. 2) the existence of the trausform function for Y,. and hence the nature of baryon settling. is indepeudenut of whether the svsteni is enibedded within a dark matter halo. aud 3) differences in initial conditions or stochastic events during the lifetimes of stellar svstenis do not invalidate the transform function."," The generality of this transform function implies that 1) the structure of one type of stellar system is not a distinct result of a process particular to that set of systems, such as violent relaxation for elliptical galaxies, but is instead the result of an interplay of all processes responsible for the generic settling of baryons in gravitational potential wells, 2) the existence of the transform function for $\Upsilon_e$, and hence the nature of baryon settling, is independent of whether the system is embedded within a dark matter halo, and 3) differences in initial conditions or stochastic events during the lifetimes of stellar systems do not invalidate the transform function."
+" By iucludiug stellar clusters. which we presuue are devoid of dark uatter. we demonstrate that the nature of the trausform ποοι for Y, depends more ou the barvon pliesies han the source of the gravitational potential."," By including stellar clusters, which we presume are devoid of dark matter, we demonstrate that the nature of the transform function for $\Upsilon_e$ depends more on the baryon physics than the source of the gravitational potential."
+ The xeseuce or abseuce of dark matter imprints uo clear οσο in the structure or kinematics of stellar clusters versus salaxies., The presence or absence of dark matter imprints no clear feature in the structure or kinematics of stellar clusters versus galaxies.
+" By iucludiug stellar clusters. which we xesune formed mostly iu single. short-lived. dramatic ‘ashiou. we demonstrate that the nature of the transform. ""nctiou for Y, is not altered radically by differences in formation hnistorv."," By including stellar clusters, which we presume formed mostly in single, short-lived, dramatic fashion, we demonstrate that the nature of the transform function for $\Upsilon_e$ is not altered radically by differences in formation history."
+ The extension of the FAL over a arecr parameter space is in reality the extension of he statement that physical processes rather than initial conditious or subsequent stochasticeveuts dominate the final outcome in the formation of stellar svsteiis., The extension of the FM over a larger parameter space is in reality the extension of the statement that physical processes rather than initial conditions or subsequent stochasticevents dominate the final outcome in the formation of stellar systems.
+We used this combined database. which is the largest ever built for such a purpose. to obtain updated estimates of the SN rates (Soc. ??)).,"We used this combined database, which is the largest ever built for such a purpose, to obtain updated estimates of the SN rates (Sec. \ref{update}) )."
+ Finally. we compared the SN rates with other SFR indicators such as the far infrared huninosity. the integrated color aud the activity of the ealaxy.," Finally, we compared the SN rates with other SFR indicators such as the far infrared luminosity, the integrated color and the activity of the galaxy."
+ Evans began lis search of SNe iu 1980 usus a 25 cur telescope., Evans began his search of SNe in 1980 using a 25 cm telescope.
+" The observations were conducted visually. which has the advantage of being very fast aud inexpensive. but the limiting maeuitude for SN discovery is not very deep Gn,=155 mae)."," The observations were conducted visually, which has the advantage of being very fast and inexpensive, but the limiting magnitude for SN discovery is not very deep $m_{\rm lim}=14.5$ mag)."
+" At the cud of 1950 the telescope was replaced with a £1 c telescope Gin,=15.0 1nag). and most recently complemented by a 100 cm telescope at Siding Spring Observatory (iin,= 16.0)."," At the end of 1985 the telescope was replaced with a 41 cm telescope $m_{\rm
+lim}= 15.0$ mag), and most recently complemented by a 100 cm telescope at Siding Spring Observatory $m_{\rm lim}=16.0$ )."
+" During the almost two decades of the search. Evaus collected over 200,000 individual observations. surveviug a sample of more than 3000 nearby galaxies."," During the almost two decades of the search, Evans collected over 200,000 individual observations, surveying a sample of more than 3000 nearby galaxies."
+ The search resulted iu the discovery of 32 SNe. and another 22 SNe. which first hac been discovered by others were also detected. for a total sample of 51 SNe.," The search resulted in the discovery of 32 SNe, and another 22 SNe, which first had been discovered by others were also detected, for a total sample of 54 SNe."
+ These umubers quality Evans’ search as the most successful amateur SN search ever a- make it very coupctitive even to professional searches., These numbers qualify Evans' search as the most successful amateur SN search ever and make it very competitive even to professional searches.
+ For the calculation of SN rates based on Ewviuis search log. we use the control time method aud the protocol described in €97.," For the calculation of SN rates based on Evans' search log, we used the control time method and the protocol described in C97."
+ The esseutial galaxy data. that is distances. morphological types. Iuninosities aud axial ratios. have been retrieved from the updated version of the RC3 catalog (de Vaucouleurs et al. 1991))," The essential galaxy data, that is distances, morphological types, luminosities and axial ratios, have been retrieved from the updated version of the RC3 catalog (de Vaucouleurs et al. \cite{rc3}) )"
+" distributed by the ""Centre de Dounéces Astronomiques de Strasbourg.", distributed by the “Centre de Donnéees Astronomiques de Strasbourg”.
+ With the improved SN statistics. the Poissonian errors on SN events were reduced aud other sources of errors became dominant. in particular in the correction of search biases.," With the improved SN statistics, the Poissonian errors on SN events were reduced and other sources of errors became dominant, in particular in the correction of search biases."
+ AILSN searches suffer from specific biases and therefore it is müaportant to compare the results of different kiuds of searches., All SN searches suffer from specific biases and therefore it is important to compare the results of different kinds of searches.
+ Because it is wnique. Evans” visual search is especially useful for couparisous with traditional photographic surveys (097).," Because it is unique, Evans' visual search is especially useful for comparisons with traditional photographic surveys (C97)."
+ Specific biases in SN searches occur because SNe appear embedded im their parcut ealaxies., Specific biases in SN searches occur because SNe appear embedded in their parent galaxies.
+ Since the detection cficiency of new objects depends on the contrast against the backeround. obviously it is more dificult to discover a SN iu the Cuiuneus) Inner regions of a galaxy than in its outskirts.," Since the detection efficiency of new objects depends on the contrast against the background, obviously it is more difficult to discover a SN in the (luminous) inner regions of a galaxy than in its outskirts."
+ This bias is at its most pronounced if one uses a wide ficld/simall scale telescope and a detector with a small dynamic ranee., This bias is at its most pronounced if one uses a wide field/small scale telescope and a detector with a small dynamic range.
+ Also the fraction of SNe ost is greater in more distaut ealaxies because of the small aneular size of the galaxy nuage., Also the fraction of SNe lost is greater in more distant galaxies because of the small angular size of the galaxy image.
+ By comparing the radial distributions of SNe iu ealaxies at differcut distances. we estimated that up to of the SNe exploding in the more distant ealaxies are lost in photographic searches using Schuüdt telescopes.," By comparing the radial distributions of SNe in galaxies at different distances, we estimated that up to of the SNe exploding in the more distant galaxies are lost in photographic searches using Schmidt telescopes."
+ The nuclear bias seeiis negligible only in CCD aud visual searches of nearby galaxies (C97)., The nuclear bias seems negligible only in CCD and visual searches of nearby galaxies (C97).
+ A precursorv exiuuinatiou of the general list of SNe reveals another severe bias against SN detection: that iu inclined. spirals., A precursory examination of the general list of SNe reveals another severe bias against SN detection: that in inclined spirals.
+ In the past. it has been claimed that this bias does not affect visual aud. CCD searches (Evaus ot al. 1989:," In the past, it has been claimed that this bias does not affect visual and CCD searches (Evans et al. \cite{ev:89};"
+ Muller ct al. 1992)), Muller et al. \cite{muller}) )
+ but so far the evidence has been incouclusive (C97)., but so far the evidence has been inconclusive (C97).
+ It is thus of iuterest to exploit the improved statistics of Evans” visual search to address this question., It is thus of interest to exploit the improved statistics of Evans' visual search to address this question.
+ We beean by calculating the overall SN rate in spiral ealaxies of different inclination both for Evans” visual search and for the combined photographic search sample constructed by C97. including the SN searches of Asiago (Cappellaro et al. 1993)).," We began by calculating the overall SN rate in spiral galaxies of different inclination both for Evans' visual search and for the combined photographic search sample constructed by C97, including the SN searches of Asiago (Cappellaro et al. \cite{PI}) ),"
+ Crimea (Tsvetkov 1983)). OCA (Pollas 1991)) aud Calaun/Tololo (IIzuuuv. ct al. 1993)).," Crimea (Tsvetkov \cite{tsv:83}) ), OCA (Pollas \cite{pol}) ) and Calánn/Tololo (Hamuy et al. \cite{mario:93}) )."
+ For this particular calculation we turned off the correction for parent Ooealaxv inclination which was included iu the recipe of C97., For this particular calculation we turned off the correction for parent galaxy inclination which was included in the recipe of C97.
+ The results are reported iu ⊺⋜⋂⋝↕↸∖↕∙∙ where the bius were chosen to give roughly the same muuber of ealaxies i cach iuclinatiou bin.," The results are reported in Table \ref{inc}, where the bins were chosen to give roughly the same number of galaxies in each inclination bin."
+ It would appear that based only on Evans’ log. the SN rate in edge-on spirals is 2.6 fold less than in face-on ones.," It would appear that based only on Evans' log, the SN rate in edge-on spirals is 2.6 fold less than in face-on ones."
+ This clearly demonstrates that the bias in inclined spirals also affects visual searches and is almost as severe as in photographic searches (2.9)., This clearly demonstrates that the bias in inclined spirals also affects visual searches and is almost as severe as in photographic searches (2.9).
+ The natural interpretation of this bias is that SNe occumne in the disk of meliued. spirals appear on the average dinuuer than those iu face-on spirals because of the increased. optical depth through the dust laver., The natural interpretation of this bias is that SNe occurring in the disk of inclined spirals appear on the average dimmer than those in face-on spirals because of the increased optical depth through the dust layer.
+ As a consequence. the probability of SN discovery in inclined spirals is reduced.," As a consequence, the probability of SN discovery in inclined spirals is reduced."
+ Iu line with this interpretation. it is expected that the bias is more severe in searches carried out in the blue baud. (e.g. photographic searches). than iu visual oues (CCD searches iu the red would be even less affected).," In line with this interpretation, it is expected that the bias is more severe in searches carried out in the blue band, (e.g. photographic searches), than in visual ones (CCD searches in the red would be even less affected)."
+ As a first order approach to correct for the inclination bias. we ean asstune a plane parallel ecometry for the distribution of dust in the disk of spirals.," As a first order approach to correct for the inclination bias, we can assume a plane parallel geometry for the distribution of dust in the disk of spirals."
+ In this case the average extinction of the SN population scales with sec;. where / is the inclination of the e:axy disk with respect to the line of sight.," In this case the average extinction of the SN population scales with $\sec i$, where $i$ is the inclination of the galaxy disk with respect to the line of sight."
+ We showed in C97 that this assuuption results in an over-correction of the SN rate iu ecec-on ealaxies., We showed in C97 that this assumption results in an over-correction of the SN rate in edge-on galaxies.
+ It was argued that this is evidence that dust is not uniformly distributed im the disk of spirals but is lustead in discrete clouds., It was argued that this is evidence that dust is not uniformly distributed in the disk of spirals but is instead in discrete clouds.
+ Until more evideuce is available we have adopted. iu analogy to COT. a conservative extinction luv that is intermediate between the sec; aud an empirical relation. derived from the assuuptiou that the SN rate is the same in face-on aud edge-on galaxies (see Section ?? for an alternative model).," Until more evidence is available we have adopted, in analogy to C97, a conservative extinction law that is intermediate between the $\sec i$ and an empirical relation, derived from the assumption that the SN rate is the same in face-on and edge-on galaxies (see Section \ref{sec_hatano} for an alternative model)."
+ When such a correction is included. we cau compute the SN rates for the complete Ένας search aud compare it with those obtained from the first teu vears of the search and with those derived from the combined photographic search suuple.," When such a correction is included, we can compute the SN rates for the complete Evans search and compare it with those obtained from the first ten years of the search and with those derived from the combined photographic search sample."
+ The results are reported in Table 2 , The results are reported in Table \ref{evares} 
+"Recent results by Gebhardtefaf(2000). ancl by Ferrarese&/Merritt.(2000) that establish a strong correlation between central black hole mass Mj, ancl the velocity dispersion σι. measured al r7./8. indicate that the central black hole is intimatelv reated to the dvnamieal structure of the galaxy.","Recent results by \citet{geb00} and by \citet{fer00} that establish a strong correlation between central black hole mass $M_{bh}$ and the velocity dispersion $\sigma_e$, measured at $r_e/8$, indicate that the central black hole is intimately related to the dynamical structure of the galaxy."
+" An earlier result by Magorrianefa£.(1998) relates Ay, linearly (ο the mass of the bulge. which suggests a similar conclusion. whie Merrifield. Forbes. Terlevieh (2000) establish a link between 3M, ancl the age of the stelar svstenm: a massive central black hole seems to grow over a period of a Gyr or so."," An earlier result by \citet{mag98} relates $M_{bh}$ linearly to the mass of the bulge, which suggests a similar conclusion, while Merrifield, Forbes, Terlevich (2000) establish a link between $M_{bh}$ and the age of the stellar system: a massive central black hole seems to grow over a period of a Gyr or so."
+ These results have already stimulated the emission of various theories that create a feed-back mechanism between bulge star formation ancl accretion of gas onto the black hole, These results have already stimulated the emission of various theories that create a feed-back mechanism between bulge star formation and accretion of gas onto the black hole
+"oulv oue carbon line has been measured we assunie NN =30 cuotethatthetotalicutralearbonmassisnotast bgtiyte Ellenseq yield,AVI uuless the MWexcitation temperature is below 208 (ulih to be the case in our targets. sce also Figure 2 : noa] ΦΩΣ.220052) so: vovaetexact choiehoien of iniis not criticaletal, for the derivationthe of nies.","only one carbon line has been measured we assume $30$ K. We note that the total neutral carbon mass is not a strong function of the assumed $T_{ex}$ unless the excitation temperature is below K (unlikely to be the case in our targets, see also Figure 2 in et 2005a) so the exact choice of is not critical for the derivation of masses."
+ For completeness. πο derwe. moleculu eas (IIs) LHASSCR assuniugalso the stancdaid factor⋅− for. starburstby environineuts shouldIvison be applicable— for SAICs— and |that QSOs Ma ο αμ πο... Atomic SolomonN WeWaκ," For completeness, we also derive molecular gas $_2$ ) masses assuming the standard conversion factor for starburst environments that should be applicable for SMGs and QSOs $\alpha_{\rm CO}$ $_{\rm
+gas}$ $L'_{\rm CO}$ $_\odot$ $^{-1}$ $^{2})^{-1}$ , Downes Solomon 1998)."
+ ∖⋅∖⊀↴↴hero asune∖ that ;We- can. he J=3 line is slelth subthermally excited (ic. ot 22005a).," We here assume that the $J$ =3 line is slightly subthermally excited (i.e., $L'_{\rm CO(3-2)}$ $\times
+L'_{\rm CO(1-0)}$, e.g., et 2005a)."
+ The CO huuinosities (not corrected for magnification) are isted iu Table 5. aud the comespoudiug masses (here corrected for. the inaguification⋅⋅. factors. given. in. Table: 1) in Table 6.," The CO luminosities (not corrected for magnification) are listed in Table 5, and the corresponding masses (here corrected for the magnification factors given in Table 4) in Table 6."
+" We note that the molecular gas masses of the subamiluueter galaxies would be higher bv ~50% if we asstuned the average Lig 5,206LLμι derived et al.. ((2010a).", We note that the molecular gas masses of the submillimeter galaxies would be higher by $\sim$ if we assumed the average $L'_{\rm CO(3-2)}$ $\times L'_{\rm CO(1-0)}$ derived by Ivison et (2010a).
+" —""m now: derive theatomicDownes carbon abundance", We can nowderive the atomic carbon abundance
+" —""m now: derive theatomicDownes carbon abundance&", We can nowderive the atomic carbon abundance
+large radii.,large radii.
+" The scatter in mass is strongly increased by matter in this region, but the orientation-induced scatter in concentration is, to a small degree, cancelled out by it."," The scatter in mass is strongly increased by matter in this region, but the orientation-induced scatter in concentration is, to a small degree, cancelled out by it."
+ Imperfect weak lensing observations smooth out part of the influence of substructure and therefore cause a slight improvement in the mass and concentration bias., Imperfect weak lensing observations smooth out part of the influence of substructure and therefore cause a slight improvement in the mass and concentration bias.
+" Shape noise, while not solely responsible for the mass and concentration scatter, contributes considerably to it, in particular for lower mass clusters."," Shape noise, while not solely responsible for the mass and concentration scatter, contributes considerably to it, in particular for lower mass clusters."
+ We have analysed a large set of mock WL observations of galaxy clusters at redshift zr~0.2 extracted from the Millennium Simulation (MS) using background galaxies with elliptical isophotes at uniform redshift zs—1.0., We have analysed a large set of mock WL observations of galaxy clusters at redshift $z_L \simeq 0.2$ extracted from the Millennium Simulation (MS) using background galaxies with elliptical isophotes at uniform redshift $z_S = 1.0$.
+ The aim of our work was to quantify the expected scatter and bias in the reconstructions of masses and concentrations of large cluster samples derived from WL surveys such as DES and LSST., The aim of our work was to quantify the expected scatter and bias in the reconstructions of masses and concentrations of large cluster samples derived from WL surveys such as DES and LSST.
+" We focus on the effect of the matter within and close to the lensing clusters, explicitly ignoring the effect of uncorrelated large-scale structure (as studied, e.g., by ?,, see below)."," We focus on the effect of the matter within and close to the lensing clusters, explicitly ignoring the effect of uncorrelated large-scale structure (as studied, e.g., by \citealt{Hoekstra_et_al_2011}, see below)."
+" Furthermore, as we are interested in the bias and scatter induced by the clusters and the weak lensing method itself, we have explicitly ignored any additional observational sources of error, such as difficulties in measuring the background galaxy ellipticities and redshifts."," Furthermore, as we are interested in the bias and scatter induced by the clusters and the weak lensing method itself, we have explicitly ignored any additional observational sources of error, such as difficulties in measuring the background galaxy ellipticities and redshifts."
+" Nevertheless, from the analysis of these mock observations, and subsequent comparison with the true cluster parameters, we can draw the following conclusions: We note that this comparison between WL-derived and ‘true’ 3D-derived concentrations is somewhat sensitive to the radial fitting range used in determining the latter; as found by ?,, a larger fitting range probing regions closer to the cluster centre tends to increase the best-fit concentration."," Nevertheless, from the analysis of these mock observations, and subsequent comparison with the true cluster parameters, we can draw the following conclusions: We note that this comparison between WL-derived and `true' 3D-derived concentrations is somewhat sensitive to the radial fitting range used in determining the latter; as found by \citet{Gao_et_al_2008}, a larger fitting range probing regions closer to the cluster centre tends to increase the best-fit concentration."
+" In this work, we have used the smallest possible truncation radius allowed by the resolution of the simulation and the mass range of our clusters (0.027209)."," In this work, we have used the smallest possible truncation radius allowed by the resolution of the simulation and the mass range of our clusters $0.02\,r_{200}$ )."
+" For an accurate calibration of the bias and scatter expected from a particular survey, it is therefore not only important to accurately model the survey, but also the way in which the theoretical comparison data is derived."," For an accurate calibration of the bias and scatter expected from a particular survey, it is therefore not only important to accurately model the survey, but also the way in which the theoretical comparison data is derived."
+" At the price of introducing a third fitting parameter, use of the Einasto profile (7) may be a possible way to remove this dependence on the radial fitting range."," At the price of introducing a third fitting parameter, use of the Einasto profile \citep{Einasto_1965} may be a possible way to remove this dependence on the radial fitting range."
+ Our analysis of the physical origin of scatter and bias suggests that the most promising way to reduce the uncertainty in concentration measurements from WL is to employ triaxial halo models as suggested by ?.., Our analysis of the physical origin of scatter and bias suggests that the most promising way to reduce the uncertainty in concentration measurements from WL is to employ triaxial halo models as suggested by \citet{Corless_King_2008}.
+" It is however important to keep in mind that not only the analysis of observational data, but also the theoretical predictions used for comparison need to use this triaxial model because the halo ‘mass’ and ‘concentration’ in a triaxial model will in general be different from their counterparts derived using spherical averaging."," It is however important to keep in mind that not only the analysis of observational data, but also the theoretical predictions used for comparison need to use this triaxial model because the halo `mass' and `concentration' in a triaxial model will in general be different from their counterparts derived using spherical averaging."
+" As stated above, in the present study we have not included the influence of uncorrelated large-scale structure at large distances from the cluster."," As stated above, in the present study we have not included the influence of uncorrelated large-scale structure at large distances from the cluster."
+" Using spherical analytic NFW haloes, ? have recently studied its effect on the"," Using spherical analytic NFW haloes, \citet{Hoekstra_et_al_2011} have recently studied its effect on the"
+increasing redshift. as interred from the downsizing of the black hole growth rate. spheroid star formation sizes both in dass and in efficiency.,"increasing redshift, as inferred from the downsizing of the black hole growth rate, spheroid star formation down-sizes both in mass and in efficiency."
+ The ratio of AGN to star formation Iuuinositv is This vields an arrow of time in the form of tracks iu ιο Magorran diagran., The ratio of AGN to star formation luminosity is This yields an arrow of time in the form of tracks in the Magorrian diagram.
+ If for cach data poiut iu the AMagorrau plane. defined by. black hole aud bulee mass. one cmn separate star formation and ACN bDnuuinosity. ren the ratio gives a vector aud hence an arrow of ine.," If for each data point in the Magorrian plane, defined by black hole and bulge mass, one can separate star formation and AGN luminosity, then the ratio gives a vector and hence an arrow of time."
+ This is of course the instantaucous trajectory of yoiuts in the Magorran plaue as viewed at any given epoch., This is of course the instantaneous trajectory of points in the Magorrian plane as viewed at any given epoch.
+ However statistically the vectors should provide ie flow of ealaxy bulecs towards the Magorrian relation., However statistically the vectors should provide the flow of galaxy bulges towards the Magorrian relation.
+ It would be interesting to construct Alagorrian flow lagraius binned over several redshift ranges., It would be interesting to construct Magorrian flow diagrams binned over several redshift ranges.
+ This would xovide Insight iuto the cosmological evolution of the flow of points in the Magorian plane., This would provide insight into the cosmological evolution of the flow of points in the Magorrian plane.
+ AGN-cnhanced star formation is mediated ly urbulent pressure aud relates spheroid star formation rate to black hole accretion rate., AGN-enhanced star formation is mediated by turbulent pressure and relates spheroid star formation rate to black hole accretion rate.
+ As tlhe ACN pressure is inereased. via jet/cocoon pressure-driven imteractious with the ISM. the induced star formation rate ds correspondingly boosted.," As the AGN pressure is increased, via jet/cocoon pressure-driven interactions with the ISM, the induced star formation rate is correspondingly boosted."
+ Downsizing for spheroid formation ids a consequence of imassive black hole downsizing., Downsizing for spheroid formation is a consequence of massive black hole downsizing.
+ The observed relation between black hole mass and spheroid velocity dispersion is obtained. with a cocfücieut (the ratio of Salpeter time to gas consuniptiou time) that provides an arrow of time.," The observed relation between black hole mass and spheroid velocity dispersion is obtained, with a coefficient (the ratio of Salpeter time to gas consumption time) that provides an arrow of time."
+ Wiehly efficient. ACN-boosted star formation is favoured at high redshift.," Highly efficient, AGN-boosted star formation is favoured at high redshift."
+ Outflows ave of order the ACN-boosted star formation rate and saturate in the super-Eddington limit., Outflows are of order the AGN-boosted star formation rate and saturate in the super-Eddington limit.
+ We eud with some relevant questions for observers that are pertinent to our model., We end with some relevant questions for observers that are pertinent to our model.
+ Is AGN activity correlated with the star formation rate?, Is AGN activity correlated with the star formation rate?
+ Was spheroid formation and black hole erowth coeval aud sviibiotic?, Was spheroid formation and black hole growth coeval and symbiotic?
+ Which came first. if either?," Which came first, if either?"
+ Is the reach of the AGN too localized to elobally affect star formation?, Is the reach of the AGN too localized to globally affect star formation?
+ Is or was feedback significant in radio-quict AGN?, Is or was feedback significant in radio-quiet AGN?
+ Are the vouugest radio sources. notably the GPC sources. templates for the earliest stages of ACN feedback. and if so. is there associated trigecring of star formation?," Are the youngest radio sources, notably the GPC sources, templates for the earliest stages of AGN feedback, and if so, is there associated triggering of star formation?"
+ If star formation is triggered bv radio cocoons. where is the evidence for a post-starburst stellar population in old radio lobes?," If star formation is triggered by radio cocoons, where is the evidence for a post-starburst stellar population in old radio lobes?"
+ if the efficient mode of star formation is due to coherent cocoon triggering as argued above. what is the evidence for spatially and temporally coordinated episodes of star formation in well-studied examples such as the Autennac?," if the efficient mode of star formation is due to coherent cocoon triggering as argued above, what is the evidence for spatially and temporally coordinated episodes of star formation in well-studied examples such as the Antennae?"
+ Tas the trigger of positive feedback in ultraluiniuous starbursts disappeared. due to a short city evele; or is it well aud truly buried im embedded AGN nuclei?," Has the trigger of positive feedback in ultraluminous starbursts disappeared, due to a short duty cycle, or is it well and truly buried in embedded AGN nuclei?"
+ Are the associated outflow rates frou AGN/starbursts of any significance for balancing the xuvouic budect of the universe. aud if so. where does the enriched debris cud wp?," Are the associated outflow rates from AGN/starbursts of any significance for balancing the baryonic budget of the universe, and if so, where does the enriched debris end up?"
+ If the eas remains in the halos of nassive ealaxies. as essentiallv all current seui-uialvtic simulations predict. why hasu't it Όσοι seen?," If the gas remains in the halos of massive galaxies, as essentially all current semi-analytic simulations predict, why hasn't it been seen?"
+ And for he modellers (who are effectively observers of the computer). how cau ACN feedback simmlatious possibly ax any predictive role if one has to begin with massive seed black holes of dubious heritage aud uucertaiu ate?," And for the modellers (who are effectively observers of the computer), how can AGN feedback simulations possibly play any predictive role if one has to begin with massive seed black holes of dubious heritage and uncertain fate?"
+ Perhaps our analytic discussion will niotivate more realistic recipes and sub-erid plivsies prescriptions for future eenerations of feedback simulations., Perhaps our analytic discussion will motivate more realistic recipes and sub-grid physics prescriptions for future generations of feedback simulations.
+ Ablauv of these questious were inspired by discussions with participants at the Oxford-Catania Workshop iu Vulcano. May. 2008. on the DIuterface. between Calaxy Formation and AGN.," Many of these questions were inspired by discussions with participants at the Oxford-Catania Workshop in Vulcano, May 2008, on the Interface between Galaxy Formation and AGN."
+ We eratefullv acknowledge all of them as well as the role of Vincenzo Antomuccio-Delogu. who organized a brilliantly topical iieetiug in the untorecttable setting of the Acolian Islauds as well as provided an opportunity for us to complete this paper.," We gratefully acknowledge all of them as well as the role of Vincenzo Antonuccio-Delogu, who organized a brilliantly topical meeting in the unforgettable setting of the Aeolian Islands as well as provided an opportunity for us to complete this paper."
+ The research of CN was funded in part by NASA eraut GO G-1730-X and a JOUΑΡΤ collaborative eraut., The research of CN was funded in part by NASA grant GO 6-1730-X and a JHU/APL collaborative grant.
+dip around 1900. as shown in Fieure 1.,"dip around 1900, as shown in Figure 1."
+ KU Cre started to show fading around 1899. aud dropped frou LL. πας (as before 1895) to ~12.0 ασ in 19021902. and then in about half a vear. it brightened to ~11.5 mae in carly 1903.," KU Cyg started to show fading around 1899, and dropped from $\sim11.4$ mag (as before 1895) to $\sim12.0$ mag in $1902-1903$, and then in about half a year, it brightened to $\sim11.5$ mag in early 1903."
+ After that. it showed fluctuations between 11.512.0 nae from 1903 to 1905.," After that, it showed fluctuations between $11.5-12.0$ mag from 1903 to 1905."
+ KU Cye is also variable ou. vear-timescales outside primary eclipses with fluctuatious 0.1 mae. as shown in the vearly binned Πο curve in the lower-rvielit paucl of Figure 1.," KU Cyg is also variable on year-timescales outside primary eclipses with fluctuations $\sim0.1$ mag, as shown in the yearly binned light curve in the lower-right panel of Figure 1."
+ Note there isa ~~0.15 mae brightening bun around 198L. which lias been discovered by Olson (1988). and is verified by DASCIL," Note there is a $\sim0.15$ mag brightening bump around 1984, which has been discovered by Olson (1988), and is verified by DASCH."
+ There secus to be a slight trend of 0.10.2 mae brighteuiug from 19109 to 1990. however. given our svstematie uncertainty over LOO vears of 0.1 mae (Taug et al.," There seems to be a slight trend of $0.1-0.2$ mag brightening from 1910 to 1990, however, given our systematic uncertainty over 100 years of $\sim0.1$ mag (Tang et al."
+ 2011). it is uot conviucing.," 2011), it is not convincing."
+ Compared with its —2000 neighbors with y=αι12 within 5 degrees. the brightening trend foi) 1910 to 1990 in KU Cre is only at the 20 level.," Compared with its $\sim 2000$ neighbors with $g=11-12$ within 5 degrees, the brightening trend from 1910 to 1990 in KU Cyg is only at the $2\sigma$ level."
+ We lave also visually checked the light curves of all these neighbors. and none of them showed similar dip around 1900 as KU Cre. which further supports that the 5-vear dip is real.," We have also visually checked the light curves of all these neighbors, and none of them showed similar dip around 1900 as KU Cyg, which further supports that the 5-year dip is real."
+ The helt curve of KU Cre. folded on the 38.139[81 days period (Olson 1988) is shown in Figure 2. with the upper panel for 1910—1990 and the lower pancl for 1890—1910.," The light curve of KU Cyg, folded on the 38.439484 days period (Olson 1988) is shown in Figure 2, with the upper panel for $1910-1990$ and the lower panel for $1890-1910$."
+ Black dots are DÀASCTI measurements. and blue arrows are upper lanits.," Black dots are DASCH measurements, and blue arrows are upper limits."
+ Both the deep primary eclipse aud thle shalow. wide secoincary eclipse can be seen in the folded light curve from 19]0 to 1990.," Both the deep primary eclipse and the shallow, wide secondary eclipse can be seen in the folded light curve from 1910 to 1990."
+ KU Cre showed larger variaülitv on uou-eclipse phases. (0.05 0.95) during 15901910 than during: 1910— 1990. with rius of 0.20 mag and 0.15 mag. resectivelv.," KU Cyg showed larger variability on non-eclipse phases $0.05-0.95$ ) during $1890-1910$ than during $1910-1990$ , with rms of 0.20 mag and 0.15 mag, respectively."
+ The 0.15 mag rns on nonm-eclie phases durius 1910.1990 further illustrates that our vpical photometric errors are below 0.15 mae., The 0.15 mag rms on non-eclipse phases during $1910-1990$ further illustrates that our typical photometric errors are below $0.15$ mag.
+ Note tha OUP coverage dving eclipse pliases (Fig., Note that our coverage during eclipse phases (Fig.
+ 2) would allow a period measurement precision of dP~P2dN/Tdp10! days for à phase uncertaintv dN~0.01 over the £7ον100 vears duration aud so it is uot possible to further refine the binary period given by Olson (1955)., 2) would allow a period measurement precision of $dP \sim P^2 dN/T \sim 4 \times 10^{-4}$ days for a phase uncertainty $dN \sim 0.01$ over the $T\sim100$ years duration and so it is not possible to further refine the binary period given by Olson (1988).
+ Since the orbital period of KU οσο is ~38.1 davs. the  Tevr facine-brightcning eveut we found from 1899 to 1901 iunust come from sources other than eclipses.," Since the orbital period of KU Cyg is $\sim38.4$ days, the $\sim5$ -yr fading-brightening event we found from 1899 to 1904 must come from sources other than eclipses."
+ Our plates are blue sensitive. where the flux is dominated bv the Ε star: This is also illustrated iu the folded heht curve. where there are up to 2.1 mae decreases caving the primary eclipses. which means the IK51ITI eiaut contributes <10% of the total flux.," Our plates are blue sensitive, where the flux is dominated by the F star; This is also illustrated in the folded light curve, where there are up to $2.4$ mag decreases during the primary eclipses, which means the K5III giant contributes $<10\%$ of the total flux."
+ Therefore. the ~0.5 mae dimming nmst be related to the F star and its accretion disk.," Therefore, the $\sim0.5$ mag dimming must be related to the F star and its accretion disk."
+ The simplest explanation of the fading-brightcning event. perhaps. is increasing disk extinction of the FE star arising from mereasing mass trauster rate.," The simplest explanation of the fading-brightening event, perhaps, is increasing disk extinction of the F star arising from increasing mass transfer rate."
+ RN Cas is an cdee-on imass-trausferiug eclipsing binary svstcm (P=32.3 days). sinilar to KU Cre.," RX Cas is an edge-on mass-transfering eclipsing binary system $P=32.3$ days), similar to KU Cyg."
+ να et al. (, Kriz et al. (
+1980) suggested iu RX Cas that when the rate of mass transfer increases. the disk density aud couscqueutly optical depth increases. leading to atteuuation of the ceutral star and resulting in fadiug in its elt curve.,"1980) suggested in RX Cas that when the rate of mass transfer increases, the disk density and consequently optical depth increases, leading to attenuation of the central star and resulting in fading in its light curve."
+ Similar effects have been observed in Be shell stars (edec-on Be stars with high esin/). where fadings are seen iu mass ejection eveuts (see e.g. Huber Floquet 1998: Catierrez-Soto ot al.," Similar effects have been observed in Be shell stars (edge-on Be stars with high $v\sin i$ ), where fadings are seen in mass ejection events (see e.g. Huber Floquet 1998; Gutierrez-Soto et al."
+ 2008)., 2008).
+" It has been found that the accretion disk surroundiug the F star iu IU Cve is thick aud dusty. with dust to eas ratio —109. and a total disk mass —105AL, (Olson 1988: Smads Plavee 1997)."," It has been found that the accretion disk surrounding the F star in KU Cyg is thick and dusty, with dust to gas ratio $\sim10^{-6}$, and a total disk mass $\sim10^{-8}-10^{-5}\ M_{\odot}$ (Olson 1988; Smak Plavec 1997)."
+ Smal Plavec (1997) estimated an extinction of 1.3 mae in V of the F star bv the disk. with 0.75 mae extinction due to the full obscuration of about half of the E star by the disk. and the remaining ~0.55 mae due to the absorption in the disk atinosphlere.," Smak Plavec (1997) estimated an extinction of 1.3 mag in V of the F star by the disk, with $\sim 0.75$ mag extinction due to the full obscuration of about half of the F star by the disk, and the remaining $\sim 0.55$ mag due to the absorption in the disk atmosphere."
+ To cause a 0.50.6 inae dimming observed in KU Cre. aud assunmüug the same gas-to-dust ratio. the disk mass excess would be on the order of the accretion disk mass. he. ~LOS10°AL...," To cause a $0.5-0.6$ mag dimming observed in KU Cyg, and assuming the same gas-to-dust ratio, the disk mass excess would be on the order of the accretion disk mass, i.e. $\sim10^{-8}-10^{-5}\ M_{\odot}$."
+ The timescale of accretion in the disk of KU Cre is a few vears (Simak Plavec 1997)., The timescale of accretion in the disk of KU Cyg is a few years (Smak Plavec 1997).
+ If we assume a ~10%?A. ejection from the I5II star arrives at the outer region of the accretion disk around 1899. during the uext few vears. the ejected mass would move inwards. so that dust erains would be levitated above the disk by radiation pressure aud plhotophoresis (sec e.g. Vinkovie 2009: Wir IInack 2009).," If we assume a $\sim10^{-8}-10^{-5}\ M_{\odot}$ ejection from the K5III star arrives at the outer region of the accretion disk around 1899, during the next few years, the ejected mass would move inwards, so that dust grains would be levitated above the disk by radiation pressure and photophoresis (see e.g. Vinkovic 2009; Wurm Haack 2009)."
+ This would result in an iucreased covering factor and a larger extinction., This would result in an increased covering factor and a larger extinction.
+ Later. when the dust moves to snaller radii in the disk aud gets closer to the FP star. if should reach he evaporation temperature and ios dust particles at high scale hcights would evaporate. leading to a xiehteuiug.," Later, when the dust moves to smaller radii in the disk and gets closer to the F star, it should reach the evaporation temperature and most dust particles at high scale heights would evaporate, leading to a brightening."
+ Aditionally. some of the dust particles may )o transported «mitwards (Takeuchi Lin 2003: Ciesla 2007). cool dowi to condensate. and disperse. leadiug ο more extinctio.," Additionally, some of the dust particles may be transported outwards (Takeuchi Lin 2003; Ciesla 2007), cool down to condensate, and disperse, leading to more extinction."
+ Moreover. when the clump of excess mass is accreted onto the F star. the energy release on the )ounudary Laver aso leads to a brightening. as discussed iu Olsou (1988) aux| Suiak Plavec (1997) to explain the (0.2 mag brightenime bump around 1981.," Moreover, when the clump of excess mass is accreted onto the F star, the energy release on the boundary layer also leads to a brightening, as discussed in Olson (1988) and Smak Plavec (1997) to explain the $\sim0.2$ mag brightening bump around 1984."
+ The above hnree processes nav together produce the fiuctuatious rom 1903 to 1905 in KU Ce., The above three processes may together produce the fluctuations from 1903 to 1905 in KU Cyg.
+ An alternative explanation is that the dinunine is caused by dust excess rather thai ass excess. i.0.. Ligher dust to gas ratio. in the disk surrounding the F star.," An alternative explanation is that the dimming is caused by dust excess rather than mass excess, i.e., higher dust to gas ratio, in the disk surrounding the F star."
+ Iu a simular wav to the above discussion. when the dust chumps move inwards aud are levitated above the disk. extinction of the F star becomes larger. aud the svsteni becomes fainter.," In a similar way to the above discussion, when the dust clumps move inwards and are levitated above the disk, extinction of the F star becomes larger, and the system becomes fainter."
+ The system brightens when the dust eraius move closer to the F star aud evaporate., The system brightens when the dust grains move closer to the F star and evaporate.
+ However. where the dust excess comes fron ds a challenging question.," However, where the dust excess comes from is a challenging question."
+ It night arise from chumps of dusty eas ojected from the W5T eiant., It might arise from clumps of dusty gas ejected from the K5III giant.
+ We have discovered i eroup of K2UT eiauts with ~1 mae dimming over 10100 vears. which do not match any known classes of variables (Tang et al.," We have discovered a group of K2III giants with $\sim1$ mag dimming over $10-100$ years, which do not match any known classes of variables (Tang et al."
+ 2010)., 2010).
+ We sugeested that the cinuuiue might be caused by dust extinction ina certain evolution stage of EK eiauts. though how the dust ects ejected is still nuclear.," We suggested that the dimming might be caused by dust extinction in a certain evolution stage of K giants, though how the dust gets ejected is still unclear."
+ The diuuuiug of KU Cre in 1599 nueht arise frou simular dust ejection processes in the ISU eiaut., The dimming of KU Cyg in 1899 might arise from similar dust ejection processes in the K5III giant.
+ The dust clumps then move to the accretion disk around the F star through Roche Lobe overflow., The dust clumps then move to the accretion disk around the F star through Roche Lobe overflow.
+ Iu smnnuux. theS-vear dip in brightness of KU Cre observed around 1900 is unique in the wav that it faded slowly aud brightened relatively quickly. which is," In summary, the5-year dip in brightness of KU Cyg observed around 1900 is unique in the way that it faded slowly and brightened relatively quickly, which is"
+A similar method to derive evidences of intrinsic absoprtion from the flux ratios in different energy ranges have been used by Alexander et al. (,A similar method to derive evidences of intrinsic absoprtion from the flux ratios in different energy ranges have been used by Alexander et al. (
+"2001, 2002) for large sample of AGNs.","2001, 2002) for large sample of AGNs."
+" However, if the sources are thick (i.e. Ny > 1034 cm’) or if the spectrum is inverted Compton(i.e. a, < 0) the ratios will not provide a good estimate of absorption In particular, for Compton thick sources, even if they are very rare among radio loud AGNs, the spectrum could be dominated by a reflection component providing ratio of ~ 5-10 indicating a low intrinsic absorption even if thea source is heavly absorbed."," However, if the sources are Compton thick (i.e. $_H$ > $^{24}$ $^{-2}$ ) or if the spectrum is inverted (i.e. $\alpha_x$ < 0) the ratios will not provide a good estimate of absorption In particular, for Compton thick sources, even if they are very rare among radio loud AGNs, the spectrum could be dominated by a reflection component providing a ratio of $\sim$ 5-10 indicating a low intrinsic absorption even if the source is heavly absorbed."
+ We note that the values of the neutral absorption estimated using the photometric method and evaluated from the spectral analysis could be different., We note that the values of the neutral absorption estimated using the photometric method and evaluated from the spectral analysis could be different.
+ This happen because the photometric method is based on the assumption that the intrinsic is a simple , This happen because the photometric method is based on the assumption that the intrinsic spectrum is a simple power-law.
+"If the has some features, as spectrumfor example emission power-law.lines, the spectrumestimates of intrinsic Ny could be different, as in the case of 3C 105, and only the detailed spectral fitting procedure will provide the correct information."," If the spectrum has some features, as for example emission lines, the estimates of intrinsic $_H$ could be different, as in the case of 3C 105, and only the detailed spectral fitting procedure will provide the correct information."
+ Much the same conclusion can be reached by the difficulty of fitting a power law plus absorption to the brighter sources., Much the same conclusion can be reached by the difficulty of fitting a power law plus absorption to the brighter sources.
+" The fact that XSPEC returns small values of a, we ascribe to the situation of a heavily absorbed spectral distribution for which it is difficult to define the actual power law above 2 keV where the effective area of the instrument is falling.", The fact that XSPEC returns small values of $\alpha_x$ we ascribe to the situation of a heavily absorbed spectral distribution for which it is difficult to define the actual power law above 2 keV where the effective area of the instrument is falling.
+" What we are dealing with is an inadequate segment of the spectrum which is essentially flat or inverted, around the peak in the spectral energy distribution."," What we are dealing with is an inadequate segment of the spectrum which is essentially flat or inverted, around the peak in the spectral energy distribution."
+" Even though we may not be able to recover the parameters of interest (a, and Ng) from the spectral fits,it is possible to demonstrate a range of intrinsic Ny column densities corresponding to some chosen range in a, by using ""fake spectra in XSPEC."," Even though we may not be able to recover the parameters of interest $\alpha_x$ and $_H$ ) from the spectral fits,it is possible to demonstrate a range of intrinsic $_H$ column densities corresponding to some chosen range in $\alpha_x$ by using 'fake' spectra in XSPEC."
+ An example for 3C 332 is shown in Fig. 5.., An example for 3C 332 is shown in Fig. \ref{fig:nh}.
+" Corresponding to the ratio of hard to medium flux of 1141, we find a range of 2 to 3.2x10?cm? for a, between 0.5 and 1."," Corresponding to the ratio of hard to medium flux of $\pm$ 1, we find a range of 2 to $\times10^{22}cm^{-2}$ for $\alpha_x$ between 0.5 and 1."
+" We have run these calculations (the “photometric method"") for each source with large hardness ratio and thus suspected to have large intrinsic absorption.", We have run these calculations (the “photometric method”) for each source with large hardness ratio and thus suspected to have large intrinsic absorption.
+ These results are given in Table 2.., These results are given in Table \ref{tab:flux}. .
+ Comparing Ny values from XSPEC to those, Comparing $_H$ values from XSPEC to those
+The four models in this paper have generally similar temperature structures throughout ihe atmosphere. as shown bv the temperature-pressure profiles in Figure 2..,"The four models in this paper have generally similar temperature structures throughout the atmosphere, as shown by the temperature-pressure profiles in Figure \ref{t_p}."
+ Although the behavior al high pressure varies between the drag and drag-Iree models. all of the moclels have (he same dav-night temperature difference at low pressures (due to the same radiative forcing set-up and very short radiative times).," Although the behavior at high pressure varies between the drag and drag-free models, all of the models have the same day-night temperature difference at low pressures (due to the same radiative forcing set-up and very short radiative times)."
+ The models with magnetic drag or extra hvperdissipation do not show the effect of small-scale numerical noise seen in the original drag-Iree model., The models with magnetic drag or extra hyperdissipation do not show the effect of small-scale numerical noise seen in the original drag-free model.
+ The temperature structure around the west terminator (0=270°) is nearly (he same as the east terminator (6=907) lor the magnetic drag models. but has a prolile that has been more altered by aclveetion in the drag-free models.," The temperature structure around the west terminator $\theta=270^\circ$ ) is nearly the same as the east terminator $\theta=90^\circ$ ) for the magnetic drag models, but has a profile that has been more altered by advection in the drag-free models."
+ The main difference between the four models is the flow structure. especially at the low pressures probed by transmission spectroscopy. where the Doppler effect could be observed.," The main difference between the four models is the flow structure, especially at the low pressures probed by transmission spectroscopy, where the Doppler effect could be observed."
+ The diflerences between the high-altitude flow in our four models can be seen in more detail in Figure 3.. which shows flow patterns across the planet at the 60 jibar level. representative of (he high-altitude regime.," The differences between the high-altitude flow in our four models can be seen in more detail in Figure \ref{fig:termu}, which shows flow patterns across the planet at the 60 $\mu$ bar level, representative of the high-altitude regime."
+ In all cases the winds are directed away [rom the substellar point towards the anti-stellar point across (he terminator., In all cases the winds are directed away from the substellar point towards the anti-stellar point across the terminator.
+ This results in a net blue-shifted wind directed. towards the observer during (transit for all models. although the magnitude of the wind speed and the details of the wind pattern vary [rom model to moclel.," This results in a net blue-shifted wind directed towards the observer during transit for all models, although the magnitude of the wind speed and the details of the wind pattern vary from model to model."
+ First we compare the no-drag models., First we compare the no-drag models.
+ Enhanced hyperdissipation in the second model reduces the numerical build-up of noise on small scales. which has the effect of making the flow more coherent. but reduces the maxinum wind speeds.," Enhanced hyperdissipation in the second drag-free model reduces the numerical build-up of noise on small scales, which has the effect of making the flow more coherent, but reduces the maximum wind speeds."
+ Nevertheless. the mean eastward wind speed across the terminator is nearly (he same for both models. and remains nearly identical between the (wo models across a wide range of pressures (5 kin tat 60 μα).," Nevertheless, the mean eastward wind speed across the terminator is nearly the same for both models, and remains nearly identical between the two models across a wide range of pressures $\sim$ 5 km $^{-1}$ at 60 $\mu$ bar)."
+ The models with magnetic drag have ΠΙΟ slower winds at this pressure., The models with magnetic drag have much slower winds at this pressure.
+ The flow in model (b) is strongest in à narrow region on either side of the terminator. which is also the area probed by transmission spectroscopy. but it has an average eastward wind speed at the terminator of only 72.3 km H*. compared to 223.8 kins H in model (a).," The flow in model (b) is strongest in a narrow region on either side of the terminator, which is also the area probed by transmission spectroscopy, but it has an average eastward wind speed at the terminator of only $\sim$ 2.3 km $^{-1}$, compared to $\sim$ 3.8 km $^{-1}$ in model (a)."
+ The magnetic drag models have similar average terminator winds throughout much of the Hatmosphere (where they have identical drag times). but at pressures less than 1 mbar the winds in the weak-drag model increase with altitude. while the winds in the strong-drag model decrease.," The magnetic drag models have similar average terminator winds throughout much of the atmosphere (where they have identical drag times), but at pressures less than 1 mbar the winds in the weak-drag model increase with altitude, while the winds in the strong-drag model decrease."
+ The transmission spectrum is obtained by dividing the spectrum obtained during transit by the stellar spectrum. thus revealing the excess absorption [rom the species that make up (he planetary atinosphere.," The transmission spectrum is obtained by dividing the spectrum obtained during transit by the stellar spectrum, thus revealing the excess absorption from the species that make up the planetary atmosphere."
+ We caleulate the attenuation in intensity of a beam of stellar light, We calculate the attenuation in intensity of a beam of stellar light
+We acknowledge the support of the Australian Research Council.,We acknowledge the support of the Australian Research Council.
+ TRW is supported by an Australian Posteraduate Award. a University of Swduev Merit Award. an Australian Astronomical Observatory PhD Scholuship aud a Denison Merit Award.," TRW is supported by an Australian Postgraduate Award, a University of Sydney Merit Award, an Australian Astronomical Observatory PhD Scholarship and a Denison Merit Award."
+The power law index p and the peak of the distribution ρω are free ,"The power law index $p$ and the peak of the distribution $\gamma_{m,o}$ are free parameters."
+"For an injection duration of /,, and total number density. nm. parameters.the energy distribution of the rate is where the hard end of the distribution σα is determined at every imestep by the size of the emitting region."," For an injection duration of $t_o$ and total number density $n_e$, the energy distribution of the rate is where the hard end of the distribution $\gamma_{\M,o}$ is determined at every timestep by the size of the emitting region."
+ IE quantities are in the frame co-moving with the llow., All quantities are in the frame co-moving with the flow.
+ ‘The distribution cools under53 and the scattering of it., The distribution cools under and the scattering of it.
+ Fitting the resulting time integrated spectrum. one may express all its characteristic requencies in terms of the corresponding energies of the ο distribution. he magnetic field D. and p.," Fitting the resulting time integrated spectrum, one may express all its characteristic frequencies in terms of the corresponding energies of the $e^-$ distribution, the magnetic field $B$, and $p$."
+ Pherefore. given the shape ofthe e— distribution ab every instant that follows from the and cooling in he magnetic and radiation fields.injection one canprescription describe heSy globalcomponent of the spectrum.," Therefore, given the shape of the $e^-$ distribution at every instant that follows from the injection prescription and cooling in the global magnetic and radiation fields, one can describe parametrically the component of the spectrum."
+ Phe same can be doneparametrically with theIC component. assuming the spectral shape results from. scattering. of monoenergetie radiation with the mean frequency ofSy.," The same can be done with the component, assuming the spectral shape results from scattering of monoenergetic radiation with the mean frequency of."
+ For the description of the time integrated spectra. the characteristic energies of thee distribution are sullicient.," For the description of the time integrated spectra, the characteristic energies of the $e^-$ distribution are sufficient."
+" In order to avoid the expensive of the evolution of the e.— distribution at every with computationself-consistent account of the7C emission. along with thetimestep. spectral distribution of the emission. the following procedure is adopted: At every imestep. the evolution of the edges 5,,and 53; of the e/— distribution is calculated: the total number density of e/— present in the region and the radiation givenenergy density and mean frequency of the raciation at all timesteps."," In order to avoid the expensive computation of the evolution of the $e^-$ distribution at every timestep, with self-consistent account of the emission, along with the spectral distribution of the emission, the following procedure is adopted: At every timestep, the evolution of the edges $\gamma_m$and $\gamma_{\M}$ of the $e^-$ distribution is calculated given the total number density of $e^-$ present in the region and the radiation energy density and mean frequency of the synchrotron radiation produced at all previous timesteps."
+ Cooling is evaluatedsvnchrotron in two cases. produced for previous in the Thomson and the extreme Wlein-Nishinaimiting (XN) accountingregime. with an scatteringadaptive timestep to insure smooth ransition.," Cooling is evaluated in two limiting cases, accounting for scattering in the Thomson and the extreme Klein-Nishina (KN) regime, with an adaptive timestep to insure smooth transition."
+" Phe law a knee which. at the end of the (assumed to last as powerlong as the developsemission). is at άρον5315,injection}."," The power law develops a knee which, at the end of the injection (assumed to last as long as the emission), is at $max\{ \gamma_{m,o}, \gamma_{\M,t_o}\}$."
+ Apart rom LTP the value +κ 15 computec were the distribution may flatten (this happens if losses in the extreme WN regime outweigh those of Sy).," Apart from $\langle\,\gamma^2 \rangle\,_{t_o}$, the value $\gamma_{\KN}$ is computed were the distribution may flatten (this happens if losses in the extreme KN regime outweigh those of )."
+ From the evolution of the of the ὁ— distribution. the total power emitted is also obtained: the edgesrelative strength of theSy and components is the timescales of ρω. 'Tving," From the evolution of the edges of the $e^-$ distribution, the total power emitted is also obtained; the relative strength of the and components is weighted by the cooling timescales of $\gamma_{m,o}$."
+ weightedthe bydescribed cooling svstem to the proposed picture of internal shocks for GRB (Mésszárros & physicalRees 1994). 1 express the needed in terms of the physical of the source. with appropriate quantitiesparameterization.," Tying the described physical system to the proposed picture of internal shocks for GRB (Mésszárros & Rees 1994), I express the needed quantities in terms of the physical properties of the source, with appropriate parameterization."
+" Internal shocks when a propertiesflow with an intrinsic variability on £i, dissipates its fluctuations.", Internal shocks develop when a flow with an intrinsic variability on $t_{var}$ dissipates its fluctuations.
+"develop The flow is further characterized by a total energy. ££. solid angle AQ. mass injection rate AZ. and duration of the activity /, Gvhieh also determine the How's bulk Lorenz factor EZ;7—ELMUI)."," The flow is further characterized by a total energy $E_o$, solid angle $\Delta \Omega$, mass injection rate $\dot{M}_o$, and duration of the activity $t_w$ (which also determine the flow's bulk Lorenz factor $\Gamma \siml \eta \simeq E_o/\dot{M}_o t_w c^2$ )."
+" Such internal shocks at à;2E7665, from the center of the activity.", Such internal shocks develop at $r_d \simeq \Gamma^2 c t_{var}$ from the center of the activity.
+ These determine the develop and. mass densities at the site to which the properties that energyare relevant to the emission processes dissipa, These determine the energy and mass densities at the dissipation site to which the properties that are relevant to the emission processes are linked.
+tionare linked. the magnetic field density is expressed as a fraction A of the Consequently.internal energy.the number energy.density of the injected ο ας a fraction ὁ of the p energy. (assuming equal numbers of p and ein the flow). and the mean energy of the injected e/— distribution 5s.," Consequently, the magnetic field energy density is expressed as a fraction $\lambda$ of the internal energy,the number density of the injected $e^-$ as a fraction $\zeta$ of the $p$ energy (assuming equal numbers of $p$ and $e^-$ in the flow), and the mean energy of the injected $e^-$ distribution by $\kappa$ ."
+" Ehe ellicicney of from the main inertia carriers to ο duringby acceleration. is passing24.=Qmm,energy with s< Qmy/m.."," The efficiency of passing energy from the main inertia carriers to $e^-$ during acceleration, is $\varepsilon_{pe} = \zeta \kappa m_e/m_p$ with $\kappa \le \zeta m_p/m_e$ ."
+" Also. in this picture. £&,= Ute."," Also, in this picture, $t_o = \Gamma t_{var}$ ."
+frame (P) and the initial spin period (Άι).,frame $\dot{P}$ ) and the initial spin period $P_{\rm init}$ ).
+" Since there is no energy injection to the pulsar now, it must be currently spinning down."," Since there is no energy injection to the pulsar now, it must be currently spinning down."
+ This implies P>0 and allows us to calculate an upper limit on the distance to the pulsar: dmax=1540 pc ," This implies $\dot{P} \geq 0$ and allows us to calculate an upper limit on the distance to the pulsar: $d_{\rm max} = 1540\,$ pc $\sigma$ )."
+"The last panel in shows Pi,>2.75 mms.", The last panel in shows $P_{\rm init} \geq 2.75$ ms.
+ We (1-σ).conclude that the pulsar in mmust have been born with a period close to the current observed value., We conclude that the pulsar in must have been born with a period close to the current observed value.
+ The birth spin period of a NS is governed by an equilibrium between the ram pressure of the accreting material and the magnetic field., The birth spin period of a NS is governed by an equilibrium between the ram pressure of the accreting material and the magnetic field.
+" NSs spun up by accretion at the Eddington rate (Meaa) are reborn as MSPs on the “rebirth line"" (Arzoumanianetal.1999).", NSs spun up by accretion at the Eddington rate $\dot{M}_{\rm edd}$ ) are reborn as MSPs on the “rebirth line” \citep{peq}.
+". shows this rebirth line on a P-P diagram, along with current positions of various pulsars from the ATNF database (Manchesteretal. 2005)."," shows this rebirth line on a $P$ $\dot{P}$ diagram, along with current positions of various pulsars from the ATNF database \citep{atnf}. ."
+. Also shown are rebirth lines for pulsars accreting at 10~!Meaq and 10-?Meaa-, Also shown are rebirth lines for pulsars accreting at $10^{-1} \dot{M}_{\rm edd}$ and $10^{-2} \dot{M}_{\rm edd}$.
+" For a pulsar radiating as a dipole, the PP product remains constant through its lifetime."," For a pulsar radiating as a dipole, the $P\dot{P}$ product remains constant through its lifetime."
+" Thus, we can calculate the initial spin-down rate of the pulsar: Pai=PP/Pinit, where Pai comes from(2)."," Thus, we can calculate the initial spin-down rate of the pulsar: $\dot{P}_{\rm init} = P \dot{P}/P_{\rm init}$, where $P_{\rm init}$ comes from."
+". If iis at ppc, then it would have been born with Pi,=3.1x10?!ss-! and B,=2.84ms."," If is at pc, then it would have been born with $\dot{P}_{\rm init} = 3.1\times 10^{-21}{\rm\,s\,s}^{-1}$ and $P_{\rm init} = 2.84\,$ ms."
+ This value is indicated with a star., This value is indicated with a star.
+" The pulsar then evolves toward the lower right, to the circle marked “E”."," The pulsar then evolves toward the lower right, to the circle marked “E”."
+" Similarly, the birth parameters for the pulsar assuming d—1.2 kpc are shown by the upward triangle — it is nearly coincident with the current parameters of the pulsar for this distance."," Similarly, the birth parameters for the pulsar assuming $d=1.2\,$ kpc are shown by the upward triangle — it is nearly coincident with the current parameters of the pulsar for this distance."
+ Other possible birth locations for the pulsar lie along the solid line passing through the star and triangle., Other possible birth locations for the pulsar lie along the solid line passing through the star and triangle.
+ The gray region shows all possible positions that ccan have occupied in its lifetime., The gray region shows all possible positions that can have occupied in its lifetime.
+" It is clear that the pulsar was born well below the rebirth line, and the mean accretion rate during the spin-up phase (the final major accretion phase) was lower than 107?Meaa-"," It is clear that the pulsar was born well below the rebirth line, and the mean accretion rate during the spin-up phase (the final major accretion phase) was lower than $10^{-2} \dot{M}_{\rm edd}$."
+ Two groups have run detailed simulations of the evolution ofJ1614-2230., Two groups have run detailed simulations of the evolution of.
+". The Linetal.(2010) model and the preferred model of Taurisetal.(2011) are qualitatively similar: the system begins as an intermediate mass X-ray binary consisting of a NS and ac4M, mmain-sequence secondary, which evolves to form a CO WD with an He envelope."," The \citet{evolve1} model and the preferred model of \citet{evolve2} are qualitatively similar: the system begins as an intermediate mass X-ray binary consisting of a NS and a $\sim 4\,$ main-sequence secondary, which evolves to form a CO WD with an He envelope."
+ The secondary accretes mass onto the NS in three phases., The secondary accretes mass onto the NS in three phases.
+" The first phase (Al) is a thermal timescale mass transfer at super-Eddington accretion rates, where the NS gains little mass."," The first phase (A1) is a thermal timescale mass transfer at super-Eddington accretion rates, where the NS gains little mass."
+" The next phase (A2) is on a nuclear timescale (—35 Myr), when the secondary is burning H in the core and envelope."," The next phase (A2) is on a nuclear timescale $\sim 35\,$ Myr), when the secondary is burning H in the core and envelope."
+" During this phase, the accretion rates are upto about a tenth of the Eddington limit."," During this phase, the accretion rates are upto about a tenth of the Eddington limit."
+" During the final accretion phase (phase AB), the secondary is burning He in its core and H in an envelope."," During the final accretion phase (phase AB), the secondary is burning He in its core and H in an envelope."
+" This causes the radius of the donor star to expand, triggering accretion at near-Eddington rates for 5 — 10MMyr."," This causes the radius of the donor star to expand, triggering accretion at near-Eddington rates for 5 – Myr."
+ The NS gains the most mass during this phase., The NS gains the most mass during this phase.
+" For typical NS parameters, accreting 0.1 — iis enough to spin them up millisecond periods (Kiziltan&Thorsett 2010)."," For typical NS parameters, accreting 0.1 – is enough to spin them up millisecond periods \citep{newage}."
+". Hence, the near-Eddington accretion in phase AB should spin the pulsar up all the way to the rebirth line."," Hence, the near-Eddington accretion in phase AB should spin the pulsar up all the way to the rebirth line."
+ This is inconsistent with our inferred birth position ofJ1614—2230., This is inconsistent with our inferred birth position of.
+". If the stellar evolution models are correct, then there is a problem with the standard formation scenario of MSPs (Radhakrishnan&Srinivasan1982;Alparetal. 1982)."," If the stellar evolution models are correct, then there is a problem with the standard formation scenario of MSPs \citep{recycling1,recycling2}."
+". In essence, the birth period depends on factors other than the magnetic field strength and accretion rates."," In essence, the birth period depends on factors other than the magnetic field strength and accretion rates."
+" For instance, Bildsten(1998,2003) proposes that accretion induces a quadrupole moment @ in the NS."," For instance, \citet{lars1,lars2} proposes that accretion induces a quadrupole moment $Q$ in the NS."
+ 'The NS then loses angular momentum by gravitational wave radiation., The NS then loses angular momentum by gravitational wave radiation.
+" Since the magnetic fields do not play any significant role in this model, the rebirth line becomes irrelevant, and mmay start its life anywhere on the P—P diagram."," Since the magnetic fields do not play any significant role in this model, the rebirth line becomes irrelevant, and may start its life anywhere on the $P$ $\dot{P}$ diagram."
+" In summary, the birth of aaway from the rebirth line requires reconsideration of angular momentum loss mechanisms."," In summary, the birth of away from the rebirth line requires reconsideration of angular momentum loss mechanisms."
+" Going forward, on-going radio observations (pulsar timing andVLBI) should improve the parallax and thereby decrease the uncertainty in the inferred P and thus better locate the pulsar in the P-P diagram."," Going forward, on-going radio observations (pulsar timing andVLBI) should improve the parallax and thereby decrease the uncertainty in the inferred $\dot P$ and thus better locate the pulsar in the $P$ $\dot P$ diagram."
+ The, The
+For a blank field observation. the umuber couuts are determined by dividing the number of siguificautlv detected sources in a sample by the area over which the sources could be detected at that level.,"For a blank field observation, the number counts are determined by dividing the number of significantly detected sources in a sample by the area over which the sources could be detected at that level."
+ In the case of a lensed feld. the sensitivity to a background source is dependent ou the position and redshift of the backeround source aud ou the redshitt of the gravitational leus.," In the case of a lensed field, the sensitivity to a background source is dependent on the position and redshift of the background source and on the redshift of the gravitational lens."
+ These lensing effects can be corrected for with a sufficicuthy: detailed mass model for the lens., These lensing effects can be corrected for with a sufficiently detailed mass model for the lens.
+ We used the models to determine the flux auplifications., We used the models to determine the flux amplifications.
+ uses miultiple-componeut lass distributious that describe the extended poteutial well of the clusters aud their more massive individual member galaxies (e.g. EKldl&neib et 11996).," uses multiple-component mass distributions that describe the extended potential well of the clusters and their more massive individual member galaxies (e.g., \markcite{kneib96}K Kneib et 1996)."
+ The mass distributions are derived from the positious of multipliiuaged features identified in high-resolution optical nuages: spectroscopic redshifts coustrain the models., The mass distributions are derived from the positions of multiply-imaged features identified in high-resolution optical images; spectroscopic redshifts constrain the models.
+ Details of the models for the A870. Ασοι and A2390 clusters cau be found oein KIznceib et ((1993). SScitz et ((1996). and I&ucib et al.," Details of the models for the A370, A851, and A2390 clusters can be found in \markcite{kneib93}K Kneib et (1993), \markcite{seitz96}S Seitz et (1996), and Kneib et al.,"
+ in preparation. respectively.," in preparation, respectively."
+ Using we mapped the background galaxies from). their observed positious back outo the source plane using kuown redshifts. where available. or to an assunied source redshift of 2=3.," Using we mapped the background galaxies from their observed positions back onto the source plane using known redshifts, where available, or to an assumed source redshift of $z=3$."
+ We eive the resulting iuplifications in column 6 of Table 3.., We give the resulting amplifications in column 6 of Table \ref{tab3}.
+ Although the amplification of a source depeuds ou both the redshift of the lens aud the redshift of the source. at the redshifts of our three cluster lenses (2=0.37 for ASTU. 2=0.1 for σοι and :=0.23 for À2390). the amplification varies ouly weakly with redshift for auv source bevond 2=Low," Although the amplification of a source depends on both the redshift of the lens and the redshift of the source, at the redshifts of our three cluster lenses $z=0.37$ for A370, $z=0.41$ for A851, and $z=0.23$ for A2390), the amplification varies only weakly with redshift for any source beyond $z=1$."
+hichamplification. BBlain ct ((1999) estimated that as long as the source redshifts are ereater than one. the svstematic uncertainties (due to the uncertainties iu both the redshift distribution of the detected sources aud the mass models of the clusters) are less than and heuce are comparable to he typical uncertainties in the absolute flux calibration of he SCUBA aps.," \markcite{blain99}B Blain et (1999) estimated that as long as the source redshifts are greater than one, the systematic uncertainties (due to the uncertainties in both the redshift distribution of the detected sources and the mass models of the clusters) are less than and hence are comparable to the typical uncertainties in the absolute flux calibration of the SCUBA maps."
+" Towever, sources with high amplifications have much arecr uucertaimties associated with positional uncertainty aud redshift incdeterminancy than sources with relatively ow alplifications."," However, sources with high amplifications have much larger uncertainties associated with positional uncertainty and redshift indeterminancy than sources with relatively low amplifications."
+ For the high amplification sources Wing rear critical lines. small Changes iu position or redshift cau result in very large variations in amplification.," For the high amplification sources lying near critical lines, small changes in position or redshift can result in very large variations in amplification."
+ In order o quantifv this effect. we computed the maxima aud nininmnm amplifications im a 3” radius region surrounding each source.," In order to quantify this effect, we computed the maximum and minimum amplifications in a $3''$ radius region surrounding each source."
+ For the sources where we do not lave a refined position from au optical identification. we give lis range iu brackets after the amplification in Table 3.," For the sources where we do not have a refined position from an optical identification, we give this range in brackets after the amplification in Table 3."
+ For sources with “typical” amplifications of 1 to 1 the effect is generally small., For sources with “typical” amplifications of 1 to 4 the effect is generally small.
+ Hoxvcever. the hieh amplification sources can casily have order of maguitude amplification uucertainties associated with their positious (aud also with their redshifts). aud this mast be allowed for iu any analysis.," However, the high amplification sources can easily have order of magnitude amplification uncertainties associated with their positions (and also with their redshifts), and this must be allowed for in any analysis."
+ Five of the sources in our sample (two in A870 aud three in A2390) fall into this category., Five of the sources in our sample (two in A370 and three in A2390) fall into this category.
+ All lie at the faint end of the sample where we cannot casily use secoucdary constraiuts. such as the lack of niultiple ages. to provide liuits on the amplificatious.," All lie at the faint end of the sample where we cannot easily use secondary constraints, such as the lack of multiple images, to provide limits on the amplifications."
+ For two of the sources in A2390 (13 and 11 in Table 3) there are no seusible upper limits to the amplifications. so we eive ouly lower οί».," For two of the sources in A2390 (13 and 14 in Table 3) there are no sensible upper limits to the amplifications, so we give only lower bounds."
+ Iu what follows we use the maxima and imininmn aiuplificatious eiveu in Table 3 to deteriuue how nuca the nuuber counts cau be changed by the amplification unicertaimties., In what follows we use the maximum and minimum amplifications given in Table 3 to determine how much the number counts can be changed by the amplification uncertainties.
+ For the direct ummber counts. we ποσα to know the source plane areas for which background galaxies would fall within the SCUBA maps aud be significantly detected. (," For the direct number counts, we need to know the source plane areas for which background galaxies would fall within the SCUBA maps and be significantly detected. ("
+Due to the expansion of the source plane. the source plane areas are suualler than the areas of the SCUBA maps.),"Due to the expansion of the source plane, the source plane areas are smaller than the areas of the SCUBA maps.)"
+ To do this. we created a erid of 2=3 backeround sources at one arcsecoud separations and used to produce the corresponding nuage plane maps with magnifications mm. (," To do this, we created a grid of $z=3$ background sources at one arcsecond separations and used to produce the corresponding image plane maps with magnifications $m$. ("
+For sources with known redshifts. we created the source plane grids at the known redshifts.),"For sources with known redshifts, we created the source plane grids at the known redshifts.)"
+ Each erid point ou the iniage plane corresponds to au area on the source plaue of 1 iresec?., Each grid point on the image plane corresponds to an area on the source plane of 1 $^2$.
+ If the sensitivities of the SCUBA maps were uniforii. we could determine the source plane area for a 20 detection at a eiven submillimeter fiux by finding the nuuber of poiuts at cach mage plane exid point that satisfied the equation 30/1f(S50jui) aud then jiultipliug that muuber by 1 arcsec?.," If the sensitivities of the SCUBA maps were uniform, we could determine the source plane area for a $3\sigma$ detection at a given submillimeter flux by finding the number of points at each image plane grid point that satisfied the equation $3\sigma/mf(850~\mu$ m) and then multiplying that number by 1 $^2$."
+" llowever. since the SCUBA maps become less seusitive towards the edges where the exposure times are less. the equation to be satisfied is instead 30,,;4/(00aar)XfF(G5U0 jn). where 0,,;, 1s the minimaun noisevt in the field and tay Is the maximum weighted exposure time."," However, since the SCUBA maps become less sensitive towards the edges where the exposure times are less, the equation to be satisfied is instead $3\sigma_{min}/(m\sqrt{t/t_{max}})<f(850~\mu$ m), where $\sigma_{min}$ is the minimum noise in the field and $t_{max}$ is the maximum weighted exposure time."
+ The uunuber of points satistvinge this equation cau then be jiultiplied by. 1 aresec?., The number of points satisfying this equation can then be multiplied by 1 $^2$.
+ Figure 39. shows the areas over which a source with a eiven flux in the source plane would be detected at the 360 level in the inage plane for each of the three cluster fields., Figure \ref{fig3} shows the areas over which a source with a given flux in the source plane would be detected at the $3\sigma$ level in the image plane for each of the three cluster fields.
+ The high maguification of the A370 cluster leus means there is less source plane area at ligh subinillimeter fluxcs in the A370 field but wich more source plane area at low submillimeter fluxes., The high magnification of the A370 cluster lens means there is less source plane area at high submillimeter fluxes in the A370 field but much more source plane area at low submillimeter fluxes.
+ The SCUBA data available ou A370 are deeper than the data on the other two clusters. which also contributes to the sizeable differences in source plane areas seen in Fie.," The SCUBA data available on A370 are deeper than the data on the other two clusters, which also contributes to the sizeable differences in source plane areas seen in Fig."
+ 39 at the fainter fluxes., \ref{fig3} at the fainter fluxes.
+smaller dispersion. of 0.08 dex.,"smaller dispersion, of 0.08 dex."
+ The origin of the dispersion is probably due to ditferences between the real N/O-O relation of the sample and the one assumed in the models., The origin of the dispersion is probably due to differences between the real N/O-O relation of the sample and the one assumed in the models.
+ We did not find any correlation of the ionization parameter with the metallicity for the objeets in our sample., We did not find any correlation of the ionization parameter with the metallicity for the objects in our sample.
+ We conclude that the combination of two line ratio predicted by photoionization models. one sensitive to the metallieity and another sensitive to the ionization parameter. which does taken into account the physical conditions of star-forming regions. gives O/H estimates close to the values derived using direct detections of electron temperatures.," We conclude that the combination of two line ratio predicted by photoionization models, one sensitive to the metallicity and another sensitive to the ionization parameter, which does taken into account the physical conditions of star-forming regions, gives O/H estimates close to the values derived using direct detections of electron temperatures."
+ This work was supported by FAPESP under grant 7., This work was supported by FAPESP under grant 2009/14787-7.
+ GH is grateful to the Spanish for support under. grant. AYA2007-67965-CO3-03 and AYA2010-21887-CO4-03. and the Comunidad de Madrid. under grant S2009/ESP-1496 (ASTROMADRID).," GH is grateful to the Spanish for support under grant AYA2007-67965-C03-03 and AYA2010-21887-C04-03, and the Comunidad de Madrid under grant S2009/ESP-1496 (ASTROMADRID)."
+ EPM is grateful to the Spanish for support under grant AYA2007-67965-CO3-02 and AYA2010-21887-CO4-02. and the Junta de a under grant TIC]|4.," EPM is grateful to the Spanish for support under grant AYA2007-67965-C03-02 and AYA2010-21887-C04-02, and the Junta de a under grant TIC114."
+The Sloan Digital Sky Survey has. discovered. Luminous quasars at redshifts as high as ο6.4. ie. when the universe was only a billion vears old.,"The Sloan Digital Sky Survey has discovered luminous quasars at redshifts as high as $z\sim6.4$, i.e., when the universe was only a billion years old."
+ The super-massive black-holes (SMDBLIID) powering these quasars have been estimated. to rave a mass of about 10 solar masses., The super-massive black-holes (SMBH) powering these quasars have been estimated to have a mass of about $10^9$ solar masses.
+ However questions regarding the galaxies that host these high recshilt quasars ive remained larecly unanswered., However questions regarding the galaxies that host these high redshift quasars have remained largely unanswered.
+ To understand. the ormation and evolution of quasars and the super-massive Xdack-holes that power them. one needs to determine several important physical parameters (such as the quasar lifetime. he ratio of black-hole mass to halo mass. and the efficiency and rate of accretion during the luminous phase). as well as the evolution of these parameters with time.," To understand the formation and evolution of quasars and the super-massive black-holes that power them, one needs to determine several important physical parameters (such as the quasar lifetime, the ratio of black-hole mass to halo mass, and the efficiency and rate of accretion during the luminous phase), as well as the evolution of these parameters with time."
+ Attempts to answer these questions generally consicer the quasar luminosity function. as it provides a tracer of the density of quasars with different luminosities as function of cosmic epoch (c.g. Llachnelt. Natarajan Rees 1998: Llaiman Loeb 1998: Ixaulfmann Lachnelt 2000: Volonteri. Eaardt Macau 2003: ντο Loch 2003)," Attempts to answer these questions generally consider the quasar luminosity function, as it provides a tracer of the density of quasars with different luminosities as function of cosmic epoch (e.g. Haehnelt, Natarajan Rees 1998; Haiman Loeb 1998; Kauffmann Haehnelt 2000; Volonteri, Haardt Madau 2003; Wyithe Loeb 2003)."
+ llowever. all these analyses are. moctel dependent ancl implicitly assume either. (a) the quasar lifetime and its evolution with redshift. and/or (b) the form. normalisation and evolution of a relation between SMDLL mass and the characteristic velocity of the host galaxy.," However, all these analyses are model dependent and implicitly assume either (a) the quasar lifetime and its evolution with redshift, and/or (b) the form, normalisation and evolution of a relation between SMBH mass and the characteristic velocity of the host galaxy."
+ The, The
+al.,al.
+ 2001) cameras were operated. in full (rame mode with 1e thin filter applied., 2001) cameras were operated in full frame mode with the thin filter applied.
+ 3oth the E864 and SGP regions overlap with the (Colless ct al., Both the F864 and SGP regions overlap with the (Colless et al.
+ 2001)., 2001).
+ The σάς is a large-scale Es»ectroscopic survey that Lully exploits the capabilities of 1e 2dE multi-fiber spectrograph on the the 3mm Anelo-Australian Telescope (AAT)., The 2dFGRS is a large-scale spectroscopic survey that fully exploits the capabilities of the 2dF multi-fiber spectrograph on the the m Anglo-Australian Telescope (AAT).
+ This is an on-going project wt aims to obtain high quality spectra and redshifts for 0000 ealaxies brighter than an extinction corrected limit of bj« l19.1mmag over an area =2000cddeg?., This is an on-going project that aims to obtain high quality spectra and redshifts for 000 galaxies brighter than an extinction corrected limit of $bj<19.4$ mag over an area $\approx 2000$ $^2$.
+ Recently. data for about 0000. &alaxies have been released (100k data release).," Recently, data for about 000 galaxies have been released (100k data release)."
+ Part. of this. unprecedented: spectroscopic database overlap with our survey and. hence optical spectroscopic information is available for galaxies in the Fs64 and SGP regions to the above magnitude lint., Part of this unprecedented spectroscopic database overlap with our survey and hence optical spectroscopic information is available for galaxies in the F864 and SGP regions to the above magnitude limit.
+ Aloreover. the central regions of both the FS64 and SCGP survey have been observed. by theROSAT satellite (Grilliths et al.," Moreover, the central regions of both the F864 and SGP survey have been observed by the satellite (Griffiths et al."
+ 1996: Georgantopoulos ct al., 1996; Georgantopoulos et al.
+ 1996) The data have been analysed: using the Science Analysis Software (SAS 5.3)., 1996) The data have been analysed using the Science Analysis Software (SAS 5.3).
+ Event files for the PN and the two MOS detectors have been produced. using the and tasks of SAS respectively., Event files for the PN and the two MOS detectors have been produced using the and tasks of SAS respectively.
+ The event files were screened for high particle background. periods by rejecting times with kkeV. count. rates higher than 25 and 15eccts/s for the PN and the two MOS. cameras respectively., The event files were screened for high particle background periods by rejecting times with keV count rates higher than 25 and cts/s for the PN and the two MOS cameras respectively.
+ These criteria excluded 5 fields (E8644. 1.5.7 and 9) from the analysis that sullered by significantly elevated and Haring particle background.," These criteria excluded 5 fields (F864–4, SGP--1, 5, 7 and 9) from the analysis that suffered by significantly elevated and flaring particle background."
+ The PN and MOS &oocd time intervals for the remaining pointings are shown in ‘Table 1., The PN and MOS good time intervals for the remaining pointings are shown in Table 1.
+ Phe differences between the PN and MOS exposure times are due to varving start and end times of individual observations., The differences between the PN and MOS exposure times are due to varying start and end times of individual observations.
+ Only events corresponding to patterns O4 for the PN and 0.12 for two MOS cameras have been kept., Only events corresponding to patterns 0–4 for the PN and 0–12 for two MOS cameras have been kept.
+ To increase the signaltonoise ratio and to reach fainter Duxes the PN and the MOS event files have been combined into à single event list using the task of SAS., To increase the signal–to–noise ratio and to reach fainter fluxes the PN and the MOS event files have been combined into a single event list using the task of SAS.
+ Images in celestial coordinates with pixel size of aaresce have been extracted in the spectral bands 2kkeV (soft) and 2-SkkeV (hard) for the merged event file., Images in celestial coordinates with pixel size of arcsec have been extracted in the spectral bands keV (soft) and keV (hard) for the merged event file.
+ Exposure maps accounting for vignetting. CCD gaps and xul pixels have been constructed for each spectral band.," Exposure maps accounting for vignetting, CCD gaps and bad pixels have been constructed for each spectral band."
+ Source detection is performed on the kkeV image using the task of CLAO with a false probability hreshold of 107 (Freeman et al., Source detection is performed on the keV image using the task of CIAO with a false probability threshold of $10^{-5}$ (Freeman et al.
+ 2002) corresponding to about 5o., 2002) corresponding to about $5\sigma$.
+ A total of 508 X-ray sources have been detected o the limit fx(0.5SkeV)zm10therestem7.," A total of 508 X-ray sources have been detected to the limit $f_X(\rm
+0.5-8\,keV)\approx10^{-14}\,erg\,s^{-1}\,cm^{-2}$."
+ The ).5-SkkeV detection lux limit. of individual. »ointings for à point source on-axis are shown in Table 1., The keV detection flux limit of individual pointings for a point source on-axis are shown in Table 1.
+ We note that the false probability threshold of 10.7 for 10 total uid. kKkeV). corresponds to a detection threshold. of about to in the soft band kkeV)., We note that the false probability threshold of $10^{-5}$ for the total band keV) corresponds to a detection threshold of about $4\sigma$ in the soft band keV).
+ This limit translates ο à typical on-axis source flux of fy(0.52keV)10erestem= in the soft kkeV. band.," This limit translates to a typical on-axis source flux of $f_X(\rm 0.5-2\,keV)\approx4 \times
+10^{-15}\,erg\,s^{-1}\,cm^{-2}$ in the soft keV band."
+ A byproduct of the source extraction. algorithm is he construction of background maps for cach spectral iud., A byproduct of the source extraction algorithm is the construction of background maps for each spectral band.
+ “Lhe mean backerouncl counts per pixel (estimated rom the background maps) are also listed in. Table 1., The mean background counts per pixel (estimated from the background maps) are also listed in Table 1.
+ A detailed. analysis of the nature of the X-ray sources detected in the Fs64 and SGP areas will be presented in a ortheoming paper (CGcorgantopoulos et al., A detailed analysis of the nature of the X-ray sources detected in the F864 and SGP areas will be presented in a forthcoming paper (Georgantopoulos et al.
+ in preparation)., in preparation).
+ To convert counts to lus the Energy Conversion Factors (ECE)of individual detectors are caleulated assuming a power law spectrum with D—17 and Galactic absorption Nyὁ10cm> appropriate for both the SCL and ES64 fields., To convert counts to flux the Energy Conversion Factors (ECF)of individual detectors are calculated assuming a power law spectrum with $\Gamma=1.7$ and Galactic absorption $N_H=2\times 10^{20} \rm {cm^{-2}}$ appropriate for both the SGP and F864 fields.
+ The mean LCL for the mosaic. of all. three electors is estimated by weighting the ECs of individual detectors by the respective exposure time., The mean ECF for the mosaic of all three detectors is estimated by weighting the ECFs of individual detectors by the respective exposure time.
+ For the encircled energy correction. accounting for the energy fraction outside the aperture within which source counts are accumulated. we adopt the calibration performed. by Cihizzardi. (2001a. 2001b).," For the encircled energy correction, accounting for the energy fraction outside the aperture within which source counts are accumulated, we adopt the calibration performed by Ghizzardi (2001a, 2001b)."
+ These studies use both PN and. MOS observations of point sources to formulate the PSE for different energies ancl oll-axis angles., These studies use both PN and MOS observations of point sources to formulate the PSF for different energies and off-axis angles.
+ In. particular. a Wing profile is fit to the data with parameters that are a function of both energv and oll-axis angle.," In particular, a King profile is fit to the data with parameters that are a function of both energy and off-axis angle."
+" The ""normal galaxy. sample used in the present study is compiled. from the LOOk data release of the 2dE Calaxy Redshift Survey.", The `normal' galaxy sample used in the present study is compiled from the 100k data release of the 2dF Galaxy Redshift Survey.
+ X. full description of the survey. design. spectroscopic observations. data reduction and. redshift measurements of the 2dEGIUS dataset. can be found. at Colless et al. (," A full description of the survey design, spectroscopic observations, data reduction and redshift measurements of the 2dFGRS dataset can be found at Colless et al. ("
+2001).,2001).
+ In wief. redshift estimation is performed. using two independent. methods: the first. cross-correlates the spectra with a library of. absorption-line galaxy and stellar templates after clipping emission lines. while the second finds and then fits emission lines.," In brief, redshift estimation is performed using two independent methods: the first cross-correlates the spectra with a library of absorption-line galaxy and stellar templates after clipping emission lines, while the second finds and then fits emission lines."
+ The automated redshift estimates are also confirmed. by visual inspection., The automated redshift estimates are also confirmed by visual inspection.
+ Occasionally. a manual recdshift is determined by itting spectral features missed by. the automated: moethod (Colless et al.," Occasionally, a manual redshift is determined by fitting spectral features missed by the automated method (Colless et al."
+ 2001)., 2001).
+ X quality Dag. Q. measuring the reliability. of. the estimated: redshifts has been assigned.," A quality flag, $Q$, measuring the reliability of the estimated redshifts has been assigned."
+ tcliable redshifts have Q=3. while Q=2 corresponds o à probable redshift and (9=1 indicates no redshift measurenent.," Reliable redshifts have $Q\ge3$, while $Q=2$ corresponds to a probable redshift and $Q=1$ indicates no redshift measurement."
+ Spectral classification of the θά galaxies has ovn performed by applying Principal Component Analysis (PCA) techniques on the observed. spectra. (Folkes et. al., Spectral classification of the 2dFGRS galaxies has been performed by applying Principal Component Analysis (PCA) techniques on the observed spectra (Folkes et al.
+ 1999: Macdgwick et al., 1999; Madgwick et al.
+ 2002)., 2002).
+ Spectral types are determined on the basis of the n parameter: where pep. pes are the projections of the first two cigenspectra derived from PCA.," Spectral types are determined on the basis of the $\eta$ parameter: where $pc_1$ , $pc_2$ are the projections of the first two eigenspectra derived from PCA."
+ Comparison. of the 4 , Comparison of the $\eta$ 
+effect we extracted the source spectrum from an annular region centred on the best source position excluding the 10 central pixels.,effect we extracted the source spectrum from an annular region centred on the best source position excluding the 10 central pixels.
+ The outer radius of the annulus was set to 40 pixels., The outer radius of the annulus was set to 40 pixels.
+ The resultant spectrum has 1201 cts., The resultant spectrum has 1201 cts.
+ An absorbed power law fits the spectrum well (y7=1.1 for 45 DOF)., An absorbed power law fits the spectrum well $\chi_\nu^2$ =1.1 for 45 DOF).
+ The value of iis compatible with the average value of the absorption along the line of sight. which indicates that IGR J19378-0617 ts a Sey 1.5 which is not intrinsically absorbed.," The value of is compatible with the average value of the absorption along the line of sight, which indicates that IGR $-$ 0617 is a Sey 1.5 which is not intrinsically absorbed."
+ IGR J232544+5842 was first reported in the third edition of the IBIS catalogue (Birdetal..2007)., IGR J23254+5842 was first reported in the third edition of the IBIS catalogue \citep{bird07}.
+. It was detected in both the 20-40 keV and 40-100 keV energy ranges. and the 18-60 significance of its detection was 6.3c out of a total of 1780 ks of observations.," It was detected in both the 20–40 keV and 40–100 keV energy ranges, and the 18–60 significance of its detection was $\sigma$ out of a total of 1780 ks of observations."
+ Within the Swift//XRT error box lies a single 2MASS source., Within the /XRT error box lies a single 2MASS source.
+ 2MASS J2352221145845307 is away from the centre of the eerror box. and it is also bright in the DSS I II images.," 2MASS J23522211+5845307 is away from the centre of the error box, and it is also bright in the DSS I II images."
+ Although a very faint source may be present in the U-filter of the UVOT data. we cannot precisely determine its magnitude.," Although a very faint source may be present in the U-filter of the UVOT data, we cannot precisely determine its magnitude."
+ We extracted a spectrum from the single Swift//XRT pointing., We extracted a spectrum from the single /XRT pointing.
+ The source is very weak with a 70 cts spectrum for a total of 5902 s exposure., The source is very weak with a 70 cts spectrum for a total of 5902 s exposure.
+ A simple power law fits the spectrum well (C-statistic value = 19.5 for 13 bins)., A simple power law fits the spectrum well (C-statistic value = 19.5 for 13 bins).
+ Adding an absorbing component improves the statistic to 16.3 for 14 bins., Adding an absorbing component improves the statistic to 16.3 for 14 bins.
+ Note that when all parameters are left free to vary they are poorly constrained., Note that when all parameters are left free to vary they are poorly constrained.
+ We therefore froze the power law photon index to 2., We therefore froze the power law photon index to 2.
+ The prefered value of mmavy indicate that the absorption is intrinsic to the source. but the poor quality of the data prevents us any firm conclusion on that matter.," The prefered value of may indicate that the absorption is intrinsic to the source, but the poor quality of the data prevents us any firm conclusion on that matter."
+ We analysed Swifi//XRT observations of 12. IGRs_ that previously lacked X-ray position at several aresee accuracy., We analysed /XRT observations of 12 IGRs that previously lacked X-ray position at several arcsec accuracy.
+ This lack of fine positions at X-ray energy either prevented a confirmation of the supposed type of the object or simply prevented nature of the object to be found., This lack of fine positions at X-ray energy either prevented a confirmation of the supposed type of the object or simply prevented nature of the object to be found.
+ The refinement of the X-ray positions allowed us to identify potential counterparts at infrared. optical and UV wavelengths for all of them.," The refinement of the X-ray positions allowed us to identify potential counterparts at infrared, optical and UV wavelengths for all of them."
+ We also report the detection of six serendipitous sources of unknown nature although in the Case of SWIFT J023405.1-322707. a KO star is the likely counterpart and thus suggests the source has a galactic origin.," We also report the detection of six serendipitous sources of unknown nature although in the case of SWIFT J023405.1+322707, a K0 star is the likely counterpart and thus suggests the source has a galactic origin."
+ All IGRs that were formerly suspected to be AGN were confirmed through our analysis as indeed being so., All IGRs that were formerly suspected to be AGN were confirmed through our analysis as indeed being so.
+ This shows that although the error box of ccan contain several candidate counterparts. when an AGN is found inside it is usually also at the origin of the hard X-ray emission.," This shows that although the error box of can contain several candidate counterparts, when an AGN is found inside it is usually also at the origin of the hard X-ray emission."
+ This is especially true for sources that have high galactic latitude (> 107))., This is especially true for sources that have high galactic latitude $>10$ ).
+ We confirm that IC 4518A is the counterpart to IGR J14579-4308. and therefore that this source is a Sey 2.," We confirm that IC 4518A is the counterpart to IGR $-$ 4308, and therefore that this source is a Sey 2."
+ We also truly identify IGR 119378—0617 as a Sey l.5 galaxy. with known infrared. radio and X-ray counterparts.," We also truly identify IGR $-$ 0617 as a Sey 1.5 galaxy, with known infrared, radio and X-ray counterparts."
+ In IGR 714579—4308 we detected a soft excess in the X-ray spectrum., In IGR $-$ 4308 we detected a soft excess in the X-ray spectrum.
+ Soft excesses have been detected in a large number of X-ray spectra of AGNs (e.g.Porquetet , Soft excesses have been detected in a large number of X-ray spectra of AGNs \citep[e.g.][]{porquet04}. .
+"The estimated luminosity of this soft excess Is 7.9x10?"" erg s7!. which is compatible with an origin intrinsic to the AGN."," The estimated luminosity of this soft excess is $\times10^{40}$ erg $^{-1}$, which is compatible with an origin intrinsic to the AGN."
+" When fitting with a nodel instead of a black body. we obtain a lower limit on the inner radius of the disk R;,>O6R;;. For the other sources. we either found in the optical and infrared surveys. faint sources within the eerror box."," When fitting with a model instead of a black body, we obtain a lower limit on the inner radius of the disk $_{in}>6$ $_{G}$ For the other sources, we either found in the optical and infrared surveys, faint sources within the error box."
+ In two cases these counterparts may be extended or a blend of sources. which prevents an identification to be given.," In two cases these counterparts may be extended or a blend of sources, which prevents an identification to be given."
+ In two cases U and UVW? counterparts were found., In two cases U and UVW2 counterparts were found.
+ For IGR J09523-6231 we first proposed a tentative Sey 2 identification., For IGR $-$ 6231 we first proposed a tentative Sey 2 identification.
+ For IGR J10147—6354 the identification as an AGN. possibly a Sey 2. seems more secure.," For IGR $-$ 6354 the identification as an AGN, possibly a Sey 2, seems more secure."
+ We. however. stress that only through optical spectroscopy of the counterpart shall the identification be firmly given.," We, however, stress that only through optical spectroscopy of the counterpart shall the identification be firmly given."
+ For the others. the fact that these objects have point sources as optical/infrared counterparts may suggest that they are Galactic sources. although this is not a definite proof.," For the others, the fact that these objects have point sources as optical/infrared counterparts may suggest that they are Galactic sources, although this is not a definite proof."
+ Only in some specific cases this is. however. strongly supported by some additional facts.," Only in some specific cases this is, however, strongly supported by some additional facts."
+ IGR J19308+0530 has an F5 star as the nost likely counterpart., IGR J19308+0530 has an F8 star as the most likely counterpart.
+ It spectrum is indicative of little intrinsic absorption (which may also suggest that it is à close object). and is very soft.," It spectrum is indicative of little intrinsic absorption (which may also suggest that it is a close object), and is very soft."
+ IGR J18490—0000 has a K-band counterpart., IGR $-$ 0000 has a K-band counterpart.
+ Its spectrum is intrinsically absorbed and resembles that of an XRB., Its spectrum is intrinsically absorbed and resembles that of an XRB.
+ Its position m the direction of the Sagittarius Arm tangent would strengthen its Galactic nature. as the arms of the Galaxy are sites with high-density of sources.," Its position in the direction of the Sagittarius Arm tangent would strengthen its Galactic nature, as the arms of the Galaxy are sites with high-density of sources."
+ Note that these sources lie at Galactic latitudes <7° wwhich may suggest that they are associated with the Galactic Plane. further supporting a Galactic origin.," Note that these sources lie at Galactic latitudes $<7$ which may suggest that they are associated with the Galactic Plane, further supporting a Galactic origin."
+ In all those cases (but IGR J19308+0530). the power law photon index returned by the spectral fit may suggest they are XRBs. although a more definite identification would require optical spectroscopy of the counterpart.and study of the temporal variability of the X-ray," In all those cases (but IGR J19308+0530), the power law photon index returned by the spectral fit may suggest they are XRBs, although a more definite identification would require optical spectroscopy of the counterpart,and study of the temporal variability of the X-ray"
+the signatures of extinction and reddening by dust.,the signatures of extinction and reddening by dust.
+ In this section. these signatures are discussed and the amount. of reddening is estimated.," In this section, these signatures are discussed and the amount of reddening is estimated."
+ We then compare the amount of reddening with the observed. X-ray absorption., We then compare the amount of reddening with the observed X-ray absorption.
+ Extinetion by cosmic dust is highly wavelength dependen and hence can change observed line [ux ratios significantly away from the intrinsic (i.e. emitted) values., Extinction by cosmic dust is highly wavelength dependent and hence can change observed line flux ratios significantly away from the intrinsic (i.e. emitted) values.
+ This provides a classic method. for determining the amount of dus extinction along the line of sight to à particular emission line region., This provides a classic method for determining the amount of dust extinction along the line of sight to a particular emission line region.
+ The relative ratios of the Balmer lines of hydrogen are often used as extinction indicators due to the fact tha they are observationally convenient (being in the optica band). strong and their intrinsic relative [ux ratios are fairly well determined from atomic theory.," The relative ratios of the Balmer lines of hydrogen are often used as extinction indicators due to the fact that they are observationally convenient (being in the optical band), strong and their intrinsic relative flux ratios are fairly well determined from atomic theory."
+ In the case of Balmer lines from AGN. one woul ideally deblend the lines into kinematically clistine components. c.g. a broad component (from the BLR) ane a narrow component (from the NLR).," In the case of Balmer lines from AGN, one would ideally deblend the lines into kinematically distinct components, e.g. a broad component (from the BLR) and a narrow component (from the NLR)."
+ One could. then obtain information about the extinction through to each component., One could then obtain information about the extinction through to each component.
+ However. high-resolution data is required. to [acilitate the deblending.," However, high-resolution data is required to facilitate the deblending."
+ To avoid introducing uncertainties due to the deblending procedure we choose to use the total Jalmer decrements: thus. the extinction. estimates below should be considered as average values over all of the emission line regions.," To avoid introducing uncertainties due to the deblending procedure we choose to use the total Balmer decrements: thus, the extinction estimates below should be considered as average values over all of the emission line regions."
+ ‘Table 2 gives the observed. Balmer decrements ancl the expected intrinsic value based upon the assumption of case-D recombination., Table 2 gives the observed Balmer decrements and the expected intrinsic value based upon the assumption of case-B recombination.
+ These decrements have been converted into the reddening. £65.V). using the standard interstellar extinction. curve of Osterbrock (1989).," These decrements have been converted into the reddening, $E(B-V)$, using the standard interstellar extinction curve of Osterbrock (1989)."
+ This. interstellar extinetion curve leacs to the expression where AK. is the observed Balmer clecrement. Ri is the intrinsic Balmer decrement and a is a constant which is given in Table 2 for the three Balmer decremoents quoted.," This interstellar extinction curve leads to the expression where ${\cal R}$ is the observed Balmer decrement, ${\cal R}_{\rm intr}$ is the intrinsic Balmer decrement and $a$ is a constant which is given in Table 2 for the three Balmer decrements quoted."
+ Table 2 also associates a hydrogen column density. Ny. with this reddening.," Table 2 also associates a hydrogen column density, $N_{\rm H}$, with this reddening."
+ This is given by (LMleiles. Ixulkani Stark 1981)," This is given by (Heiles, Kulkani Stark 1981)"
+similar progenitor model.,similar progenitor model.
+" Some general features of their results are: i) the peak luminosity reaches ~10*° erg/s, the duration of the emission is a few minutes, iii) ii)the radiation temperature is of the order of 109 K. Despite the approximate treatment of emissions from the shock, light curves calculated by our method successfully reproduce the features."," Some general features of their results are: i) the peak luminosity reaches $\sim 10^{45}$ erg/s, ii) the duration of the emission is a few minutes, iii) the radiation temperature is of the order of $10^6$ K. Despite the approximate treatment of emissions from the shock, light curves calculated by our method successfully reproduce the features."
+ Figure 3 shows light curves for various viewing angles O and a fixed a(=0.5)., Figure \ref{view} shows light curves for various viewing angles $\Theta$ and a fixed $\alpha(=0.5)$.
+ A clear distinction between models with small viewing angles and those with large viewing angles is recognized., A clear distinction between models with small viewing angles and those with large viewing angles is recognized.
+" For small viewing angle models, the luminosity gradually rises after a sudden brightening and eventually reaches the peak value."," For small viewing angle models, the luminosity gradually rises after a sudden brightening and eventually reaches the peak value."
+" For large viewing angle models, on the other hand, the luminosity suddenly reaches the peak value and then begins to decline."," For large viewing angle models, on the other hand, the luminosity suddenly reaches the peak value and then begins to decline."
+ This behavior is easily explained by considering how photons emitted from the stellar surface travel to the observer., This behavior is easily explained by considering how photons emitted from the stellar surface travel to the observer.
+" In the case of small viewing angles, the shock wave initially emerges from the point on the symmetry axis (00~ 0) and then points with 09>0 experience the shock emergence one after another."," In the case of small viewing angles, the shock wave initially emerges from the point on the symmetry axis $\theta_0\simeq0$ ) and then points with $\theta_0>0$ experience the shock emergence one after another."
+"Because the emergence time distributions Τα,0) flattens at 0=7/2 as shown in Figure 1,, the shock emergence at points with 0o~7/2 is expected to occur almost simultaneously.","Because the emergence time distributions $T(\alpha,\theta)$ flattens at $\theta=\pi/2$ as shown in Figure \ref{emergence}, the shock emergence at points with $\theta_0\simeq \pi/2$ is expected to occur almost simultaneously."
+" In other words, in a later epoch, the observer receives more photons than in the earlier epoch."," In other words, in a later epoch, the observer receives more photons than in the earlier epoch."
+ This is the reason why the luminosity reaches the peak value after a gradual increase for small viewing angle., This is the reason why the luminosity reaches the peak value after a gradual increase for small viewing angle.
+" Next, we consider models with large viewing angles."," Next, we consider models with large viewing angles."
+" In this case, the shock wave initially emerges from points with 69~7/2."," In this case, the shock wave initially emerges from points with $\theta_0\simeq\pi/2$."
+ Photons emitted from the points need longer time to reach the position of the observer than those from points with 69«2/2., Photons emitted from the points need longer time to reach the position of the observer than those from points with $\theta_0<\pi/2$.
+" Because of the combined effect of the time delay and the early emergence of the shock, photons from points with 09~7/2 reach the position of the observer almost simultaneously, which leads to a more sudden brightening to the peak Iuminosity."," Because of the combined effect of the time delay and the early emergence of the shock, photons from points with $\theta_0\simeq \pi/2$ reach the position of the observer almost simultaneously, which leads to a more sudden brightening to the peak luminosity."
+" Next, we investigate how the degree of asphericity of an explosion, which is characterized bythe parameter a in eq. (1)),"," Next, we investigate how the degree of asphericity of an explosion, which is characterized bythe parameter $\alpha$ in eq. \ref{initial}) ),"
+ affects the light curve during the shock breakout., affects the light curve during the shock breakout.
+ Figure 4 shows light curves for a fixed viewing angle O(-0°) and various a., Figure \ref{f4} shows light curves for a fixed viewing angle $\Theta(=0\degr)$ and various $\alpha$.
+" The dash-dotted line represents the model with a=0, i.e., a spherical explosion."," The dash-dotted line represents the model with $\alpha=0$, i.e., a spherical explosion."
+" In this case, the shock breakout simultaneously occurs in the same way at every points on the stellar surface."," In this case, the shock breakout simultaneously occurs in the same way at every points on the stellar surface."
+" Therefore, the luminosity does not evolve with time."," Therefore, the luminosity does not evolve with time."
+" The duration of the emission is roughly determined by light-travel-time across the stellar radius, R/c~100 s. For models with a larger a, the duration of the emission is longer than that of the model with a= 0."," The duration of the emission is roughly determined by light-travel-time across the stellar radius, $R/c\simeq 100$ s. For models with a larger $\alpha$ , the duration of the emission is longer than that of the model with $\alpha=0$ ."
+ This is, This is
+Galaxy clusters are dvuamic systems m which tens to hundreds of galaxies are plowing through ho intracluster eas.,Galaxy clusters are dynamic systems in which tens to hundreds of galaxies are plowing through hot intracluster gas.
+ In the core of the cluster. tlre| galaxy deusity is the highest. aud star forming eaOs:axies preferentially avoid these regions (?77)..," In the core of the cluster, the galaxy density is the highest, and star forming galaxies preferentially avoid these regions \citep{ken83,dre99,gav09}."
+ However. at the outskirts of the cluster many eaaxies show starburst activity detected at Mid-I1raved (IR) (27) and radio (2?) wavelengths. as weIL as in optical eiuission lines (?)..," However, at the outskirts of the cluster many galaxies show starburst activity detected at Mid-Infrared (IR) \citep{koy10,gal09} and radio \citep{hec85} wavelengths, as well as in optical emission lines \citep{rin05}."
+ Coma. at z~0.023. is our nearest rich galaxy ster.," Coma, at $\sim$ 0.023, is our nearest rich galaxy cluster."
+ As such. it allows for a level of detailed y.dy currently unavailable for more cistant sters.," As such, it allows for a level of detailed study currently unavailable for more distant clusters."
+ ? studied the R-haucd hIuuünositv amction. finding no difference in the evolutionary istorv of galaxies at the outskirts compared to AC5e at the cluster core.," \citet{mob03} studied the R-band luminosity function, finding no difference in the evolutionary history of galaxies at the outskirts compared to those at the cluster core."
+ However. the study was based on a spectroscopic sample aud less deep than other studies.," However, the study was based on a spectroscopic sample and less deep than other studies."
+ More recent studies. however. have reveaed a segregation of galaxy activity with local deusivi Calasies at the cluster ceuter are oldest (?1).. with the fraction of blue. spiral. (??7).. aud «wart (77) ealaxies Increasing with distauce TOM the core.," More recent studies, however, have revealed a segregation of galaxy activity with local density: Galaxies at the cluster center are oldest \citep{pog01,smi08}, with the fraction of blue, spiral, \citep{ter01,agu04,biv02}, and dwarf \citep{ode02,gav10} galaxies increasing with distance from the core."
+ Iu the radio. à Iuniuositv functio1 las beon well studied down to a very low surface ariel:nesses (7)..," In the radio, a luminosity function has been well studied down to a very low surface brightnesses \citep{mil09}."
+ These faint radio galaxies are vast dἩ muuber. have high specific star formatio- rates (sSFRs). strong optical lines. aud appear to inve nudergone a recent burst of star formation. lor totrein quenching.," These faint radio galaxies are vast in number, have high specific star formation rates (sSFRs), strong optical lines, and appear to have undergone a recent burst of star formation, prior to their quenching."
+ Moetallicitv eradieuts m cluster eaaxies have also been noted. aud liuked Q the age segregation. as well as to pressure confu10110it from the deep poteutial at the cluster core (?7)..," Metallicity gradients in cluster galaxies have also been noted, and linked to the age segregation, as well as to pressure confinement from the deep potential at the cluster core \citep{pog01,car02}."
+ A scenario in which the galaxy activifv Is quenched during iufall toward the dense cster core ds supported by ealaxy aliguinenuts MOM," A scenario in which the galaxy activity is quenched during infall toward the dense cluster core is supported by galaxy alignments \citep{ada09}, ,"
+"We present a set of five simulations, each of which uses the initial conditions described above, the only difference being the opacity used.","We present a set of five simulations, each of which uses the initial conditions described above, the only difference being the opacity used."
+" Simulation (A, B, C, D, E) has an opacity table that is scaled by a constant value of (1/10, 1/3, 1, 3, 10)."," Simulation (A, B, C, D, E) has an opacity table that is scaled by a constant value of (1/10, 1/3, 1, 3, 10)."
+" Thus, Simulation C has an estimated physical opacity for solar metallicity, while Simulations A and B have reduced opacities, and Simulations D and E have increased opacities."," Thus, Simulation C has an estimated physical opacity for solar metallicity, while Simulations A and B have reduced opacities, and Simulations D and E have increased opacities."
+" Physical changes in opacity could be the result of grain growth (?),, grain evolution via the passage of spiral arms (?),, or formation in an environment with a non-solar metallicity."," Physical changes in opacity could be the result of grain growth \citep*{Birnstiel2010}, grain evolution via the passage of spiral arms \citep*{Podolak2011}, or formation in an environment with a non-solar metallicity."
+" Our goal, however, is not to reproduce different physical environments, but rather to explore the necessary conditions for gravitational fragmentation."," Our goal, however, is not to reproduce different physical environments, but rather to explore the necessary conditions for gravitational fragmentation."
+" In this context, a simple scaling of the opacity table is both sufficient and desirable."," In this context, a simple scaling of the opacity table is both sufficient and desirable."
+ All five simulations evolved in a similar fashion over the first 2.5 ORPs., All five simulations evolved in a similar fashion over the first 2.5 ORPs.
+ High mode-number spiral structure developed slowly from SPH Poisson noise in each of the discs over this time until settling down to a transitioning state of two or three spiral arms., High mode-number spiral structure developed slowly from SPH Poisson noise in each of the discs over this time until settling down to a transitioning state of two or three spiral arms.
+" The transition from axisymmetric initial condition to spiral structure is observed to be smooth, with no strong transients."," The transition from axisymmetric initial condition to spiral structure is observed to be smooth, with no strong transients."
+ The final states of the simulations are shown in Figure 2.., The final states of the simulations are shown in Figure \ref{figDiscEndState}.
+" Simulations C, D, and E have been evolved for roughly 8.5 ORPs without fragmentation having taken place (although strong spiral two-arm over-densities may persist), while Simulation A has fragmented with two objects forming, and Simulation B has fragmented with one object forming."," Simulations C, D, and E have been evolved for roughly 8.5 ORPs without fragmentation having taken place (although strong spiral two-arm over-densities may persist), while Simulation A has fragmented with two objects forming, and Simulation B has fragmented with one object forming."
+" This set of simulations, therefore, demonstrates a transition from non-fragmentation to fragmentation, as a function of the opacity scaling."," This set of simulations, therefore, demonstrates a transition from non-fragmentation to fragmentation, as a function of the opacity scaling."
+" For a patch of an optically thick disk, the cooling time is approximately where & is the opacity (?).."," For a patch of an optically thick disk, the cooling time is approximately where $\kappa$ is the opacity \citep{Rafikov2007}."
+" Hence, the cooling time is directly proportional to the opacity, and our set of simulations offers a means of exploring the fragmentation boundary as a function of cooling time in a manner similar to the simulations of ?.."," Hence, the cooling time is directly proportional to the opacity, and our set of simulations offers a means of exploring the fragmentation boundary as a function of cooling time in a manner similar to the simulations of \citet{Gammie2001}."
+" The difference is that our simulations use realistic radiative cooling (even though the opacities may be scaled), rather than (-prescription cooling."," The difference is that our simulations use realistic radiative cooling (even though the opacities may be scaled), rather than $\beta$ -prescription cooling."
+ We know from the cooling criterion that reducing the cooling time (by reducing the opacity) will eventually lead to fragmentation., We know from the cooling criterion that reducing the cooling time (by reducing the opacity) will eventually lead to fragmentation.
+" Thus, that Simulations A and B fragment is consistent with this picture."," Thus, that Simulations A and B fragment is consistent with this picture."
+" However, it is possible to use these simulations to better understand why exactly fragmentation takes place."," However, it is possible to use these simulations to better understand why exactly fragmentation takes place."
+" Figure 3 shows the five simulations at roughly the same time, shortly before fragmentation took place in Simulations A and B. The difference in the structure of the five discs at this time offers evidence of a detailed description for why Simulations A and B fragment, while Simulations C, D, and E do not."," Figure \ref{figDiscMidState} shows the five simulations at roughly the same time, shortly before fragmentation took place in Simulations A and B. The difference in the structure of the five discs at this time offers evidence of a detailed description for why Simulations A and B fragment, while Simulations C, D, and E do not."
+" As can be observed, as the cooling time decreases (as the opacity decreases), spiral arms in a disc become thinner and more over-dense and this makes fragmentation more likely to occur."," As can be observed, as the cooling time decreases (as the opacity decreases), spiral arms in a disc become thinner and more over-dense and this makes fragmentation more likely to occur."
+"PR, lic on the Fundameutal Plaine of ETCs (c.e.. Dernaudi et al","$R_e$ lie on the Fundamental Plane of ETGs (e.g., Bernardi et al."
+ 2003)., 2003).
+" The free parameters defining the dark matter were chosen in agreement with the results from dynamical modeling of the observed motions of stars. planetary nebulae and elobular chisters at simall aud large radi: these indicate that the dark matter beeins to be dynamically iuportaut at 3A. ίοιο, Saglia ct al."," The free parameters defining the dark matter were chosen in agreement with the results from dynamical modeling of the observed motions of stars, planetary nebulae and globular clusters at small and large radii; these indicate that the dark matter begins to be dynamically important at $R_e$ (e.g., Saglia et al."
+ 1992. Cappellari et al.," 1992, Cappellari et al."
+ 2006. Weijmanus et al.," 2006, Weijmans et al."
+ 2009. Shen Gebhardt 2010).," 2009, Shen Gebhardt 2010)."
+" This requires that §=1,/R,>1. and R=AL,(AL.3 or 5 (the latter value corresponding to the barvou-to-total nass ratio of WALAP. Isomatsu et al."," This requires that $\beta =r_h/R_e>1$, and ${\cal R}=M_h/M_*=3$ or 5 (the latter value corresponding to the baryon-to-total mass ratio of WMAP, Komatsu et al."
+ 2009)., 2009).
+" By solving nuucrically the Jeans equations for the three mass components in the isotropic orbits case (e.@.. Binney Tremaine 1987). these choices produce AL. Ag, aud νο (the diark-to-Inuinous mass ratio within R.)."," By solving numerically the Jeans equations for the three mass components in the isotropic orbits case (e.g., Binney Tremaine 1987), these choices produce $M_*$, $M_h$ and ${\cal R}_e$ (the dark-to-luminous mass ratio within $R_e$ )."
+ Reasonable values of AL/Lp=flI0)M.fLp... and νι=(0.2.1.0 are obtained.," Reasonable values of $M_*/L_B=(4-10) \, M_{\odot}/L_{B,\odot}$, and ${\cal R}_e =
+0.2-1.0$ are obtained."
+ The main properties of a few representative mass models are shown in Fig. 1.., The main properties of a few representative mass models are shown in Fig. \ref{f1}.
+" For a consistent comparison between observed 27s aud he characteristic temperatures derived for the mass uodels. the ceutral stellar velocity dispersion 6,. mist 0 the sane for observed. ETCs aud models."," For a consistent comparison between observed $T$ 's and the characteristic temperatures derived for the mass models, the central stellar velocity dispersion $\sigma_c$ must be the same for observed ETGs and models."
+ Typically. or nearby well observed ETCs. the value of σι. is that of the projected aud liminosity-weighted average within an aperture of radius ιδ.," Typically, for nearby well observed ETGs, the value of $\sigma_c$ is that of the projected and luminosity-weighted average within an aperture of radius $R_e/8$."
+" Therefore. when defining a uass niodel. the chosen value of o, was assigued to this quantity."," Therefore, when defining a mass model, the chosen value of $\sigma_c$ was assigned to this quantity."
+ Finally. streiuniug motions as stellar rotation are not cousidered in these models (possible heating frou hese motions is discussed by Ciotti Pellegrini 1996).," Finally, streaming motions as stellar rotation are not considered in these models (possible heating from these motions is discussed by Ciotti Pellegrini 1996)."
+ We investigate here. the relationship between the observed L's and those expected from the various sources of heating (stellar motions. gravitational potoeutial. SNlas). or during the escape of the hot eas.," We investigate here the relationship between the observed $T$ 's and those expected from the various sources of heating (stellar motions, gravitational potential, SNIa's), or during the escape of the hot gas."
+ For this purpose. Fig.," For this purpose, Fig."
+" 2. shows the run with o, of the various temperatures defined in Sects. ο νεα εν,"," \ref{f2} shows the run with $\sigma_c$ of the various temperatures defined in Sects. \ref{inj}, \ref{egp},"
+ and ??.. together with the distribution of the observed T values from BNF (Sect. ??)):," and \ref{esc}, together with the distribution of the observed $T$ values from BKF (Sect. \ref{intro}) );"
+" for the BISF sample. the value of 0, is the huninosity weighted average within A,./8. aud has Όσσα taken from SAURON studies for 12 ETCs (Iuutschuer et al."," for the BKF sample, the value of $\sigma_c$ is the luminosity weighted average within $R_e/8$, and has been taken from SAURON studies for 12 ETGs (Kuntschner et al."
+ 2010). for the remaining cases from the references in the Uvperleda catalog (see Tab. 1)).," 2010), for the remaining cases from the references in the Hyperleda catalog (see Tab. \ref{tab1}) )."
+ The temperatures defined in Sects. ?7.. 77..," The temperatures defined in Sects. \ref{inj}, \ref{egp},"
+ and ?? are inass-weielted averages. which is required when discussing energetic aspects of the eas (e... the energy required for escape as 1ieastred bv <Toya2: compared with the input energv from SNIaπι," and \ref{esc} are mass-weighted averages, which is required when discussing energetic aspects of the gas (e.g., the energy required for escape as measured by $<T_{\rm grav}^->$ compared with the input energy from SNIa's)."
+ When a direct conrparison is made with observed Zs. it nimmst be noted that the latter coincide with mass-weighted averages oulv if the ISA las everswhere oue single temperature value: if the eas is multi-phase. or its temperature profile has a eracicut. a single Z7 value measured from the spectrum of the iuteerated cussion will be close to an Cluission-weighted average (0.9. Ciotti Pellegrini 2008. im 2011).," When a direct comparison is made with observed $T$ 's, it must be noted that the latter coincide with mass-weighted averages only if the ISM has everywhere one single temperature value; if the gas is multi-phase, or its temperature profile has a gradient, a single $T$ value measured from the spectrum of the integrated emission will be close to an emission-weighted average (e.g., Ciotti Pellegrini 2008, Kim 2011)."
+ This means that. since the densest reeion is thecentral one. the measured Fs τοπ to be closer to the central values than the mass-woeiehted ones.," This means that, since the densest region is thecentral one, the measured $T$ 's tend to be closer to the central values than the mass-weighted ones."
+ The temperature profiles observed with Chandra change continuously in shape. as the euission-weighted average T decreases from <1 keV to ~0.3 keV: they switch from a flat ceutral profile that iucreases outward of 0.5R.. to a quasi-isotherimal profile. to a profile with a negative eradieut (Diehl Statler 2008. Nagino Matsushita 2009).," The temperature profiles observed with $Chandra$ change continuously in shape, as the emission-weighted average $T$ decreases from $\lsim 1$ keV to $\sim 0.3 $ keV: they switch from a flat central profile that increases outward of $\sim 0.5 R_e$, to a quasi-isothermal profile, to a profile with a negative gradient (Diehl Statler 2008, Nagino Matsushita 2009)."
+ Therefore. the lowest observed T's prestunably associated with the last category. may be luwger than mass-weighted values. intermediate Ts may be the closest to mass-weighted averages. while the largest observed Pos imay be lower fhan ias averages.," Therefore, the lowest observed $T$ 's, presumably associated with the last category, may be larger than mass-weighted values, intermediate $T$ 's may be the closest to mass-weighted averages, while the largest observed $T$ 's may be lower than mass-weighted averages."
+" Another aspect to recall is that the temperatures defined in Sects. 27.. 27.,"," Another aspect to recall is that the temperatures defined in Sects. \ref{inj}, \ref{egp},"
+" and 2? refer to a gas distribution with ωςXp. this is appropriate for the continuously injected ο for <<Tj,;>). while it may be less accurate when comparing observed Ts with <Te"">. or when discussing the cucrectics of the whole gas content of an ETGs by means of <Tl> aud <Dy>. neo the bulk of the hot ISMD may have a differcut distribution from the stars."," and \ref{esc} refer to a gas distribution with $\rho_{gas} \propto
+ \rho_*$; this is appropriate for the continuously injected gas (e.g., for $<T_{inj}>$ ), while it may be less accurate when comparing observed $T$ 's with $<T_{\rm esc}^{sub}>$, or when discussing the energetics of the whole gas content of an ETGs by means of $<T_{\rm grav}^+>$ and $<T_{\rm grav}^->$, since the bulk of the hot ISM may have a different distribution from the stars."
+" For example. the observed X-rav brightuess profile of gas-rich ETCs was found to follow the optical onc. which was taken as evidence that roughly pyasXp, (eg. Sarazin White 1988. Fabbiauuo 1989)."," For example, the observed X-ray brightness profile of gas-rich ETGs was found to follow the optical one, which was taken as evidence that roughly $\rho_{gas}\propto {\sqrt \rho_*}$ (e.g., Sarazin White 1988, Fabbiano 1989)."
+ For eas-poor ETCs tosting galactic winds. the modeling shows that the xofile pqus(r) will again be shallower thau ptr). though rot as uuuch as in the previous case (see. eg. White Chevalier 1983).," For gas-poor ETGs hosting galactic winds, the modeling shows that the profile $\rho_{gas} (r)$ will again be shallower than $\rho_* (r)$, though not as much as in the previous case (see, e.g., White Chevalier 1983)."
+ TE pyas has a flatter radial profile han p.. then it is casy to show that its inass-weielited <Ty> will be larger than derived using Eq. L.," If $\rho_{gas}$ has a flatter radial profile than $\rho_*$, then it is easy to show that its mass-weighted $<T_{\rm grav}^+>$ will be larger than derived using Eq. \ref{eq:lgp},"
+" aud its massweighted <Ty.> aud <TA""?> will he ower than derived using Eqs."," and its mass-weighted $<T_{\rm
+ grav}^->$ and $<T_{\rm esc}^{sub}>$ will be lower than derived using Eqs."
+ 6 anc 7.., \ref{eq:lm} and \ref{eq:tent}.
+" Iu conclusion. he comparison of observed T's with imass-weightecd expectations is the best that can be done curreuth, in a eeneral analysis as that of the preseut work. though witli he warnings above."," In conclusion, the comparison of observed $T$ 's with mass-weighted expectations is the best that can be done currently, in a general analysis as that of the present work, though with the warnings above."
+" Note. however. that all arguineuts and couclusious below remain valid or are strenghtcued. when taking into account the above considerations about observed P's. or about the modificatious to <<Tiloebay>. zTION> aud cTals,[Sd"," Note, however, that all arguments and conclusions below remain valid or are strenghtened, when taking into account the above considerations about observed $T$ 's, or about the modifications to $<T_{\rm
+ grav}^+>$, $<T_{\rm grav}^->$ and $<T_{\rm esc}^{sub}>$."
+ Iun the left panel of Fig., In the left panel of Fig.
+" 2. the observed Ts are conrpared with approximate estimates of the stellar temperature T5. of Πω). aud of the escape temperature IT, (Sect. 22))."," \ref{f2} the observed $T$ 's are compared with approximate estimates of the stellar temperature $T_{\sigma}$, of $T_{inj}$, and of the escape temperature $4T_{\sigma}$ (Sect. \ref{esc}) )."
+ The gas Iuniuositv is also iudicated witli different colors. having erouped the Ly values in three ranges. choscu to have a roughly equal uuuber of ETCs in cach range.," The gas luminosity is also indicated with different colors, having grouped the $L_X$ values in three ranges, chosen to have a roughly equal number of ETGs in each range."
+ This erouping gives au indication of the eas flow status. based on previous works: a galactic wiud leaving the galaxy with a supersonic velocity has Ly<107 ore | (ee. Mathews Baker 1971. Trinchieri et al.," This grouping gives an indication of the gas flow status, based on previous works: a galactic wind leaving the galaxy with a supersonic velocity has $L_X<10^{38}$ erg $^{-1}$ (e.g., Mathews Baker 1971, Trinchieri et al."
+" 2008). global subsomic outflows aud partial winds can reach Ly~10"" eres 1 (Ciotti et al."," 2008), global subsonic outflows and partial winds can reach $L_X\sim 10^{40}$ erg $^{-1}$ (Ciotti et al."
+ 1991. Pellegrini Ciotti 1998). and a central inflow becomes increasingly more imuportaunt in ETCGs of increasingly larecr Ly.," 1991, Pellegrini Ciotti 1998), and a central inflow becomes increasingly more important in ETGs of increasingly larger $L_X$ ."
+ Magenta ETCs (107 eres 1Lyc15«107 eres 1) should theu host winds. subsonic outflows. aud partial winds with a very small inflowine region of radius <100 per evan ETGs (1.5.LO? ere beEy<12.10 erg 1) should host subsouic outflows and partial winds with an increasingly larger iufiowiug region (of radius up to a few hundreds pe): black ETCs are hot gas-rich aud," Magenta ETGs $10^{38}$ erg $^{-1}<L_X<1.5\times 10^{39}$ erg $^{-1}$ ) should then host winds, subsonic outflows, and partial winds with a very small inflowing region of radius $<100$ pc; cyan ETGs $1.5\times
+10^{39}$ erg $^{-1}<L_X<1.2\times 10^{40}$ erg $^{-1}$ ) should host subsonic outflows and partial winds with an increasingly larger inflowing region (of radius up to a few hundreds pc); black ETGs are hot gas-rich and"
+" where the tilde over A!4,» in the above expressions is introduced. to. indicate. explicit factoring out of the eravitational wave amplitude /.","= ], where the tilde over $R_{+1,+2}$ in the above expressions is introduced to indicate explicit factoring out of the gravitational wave amplitude $h$."
+ Usually owe know only statistical properties οἱ the gravitational wave field., Usually we know only statistical properties of the gravitational wave field.
+" Stationary statistically. homogeneous and isotropic gravitational wave field possesses the following properties: hihi)> = hh) bu) ""where the brackets denote ensemble naveraging over all possible realizations. anc {ο is the metric power spectrum per logarithmic interval of &."," Stationary statistically homogeneous and isotropic gravitational wave field possesses the following properties: <h_s(k^i)> = <h_s(k^i) (k'^i)> = ^3(k^i-k'^i), where the brackets denote ensemble averaging over all possible realizations, and $P_h(k)$ is the metric power spectrum per logarithmic interval of $k$."
+ Using (15)). we can calculate the statistical properties of the corresponding shifts in frequency derivatives ο) and ο).," Using \ref{gwstatprop}) ), we can calculate the statistical properties of the corresponding shifts in frequency derivatives $\dot{\nu}(t)$ and $\ddot{\nu}(t)$."
+" Using (11)) and (15)). and taking into account the orthogonality property (9)). after straight Forward. calculations. we arrive at the following statistical expressions: = 0 =f HED,URE lett (CUR) where we have mintroduced the SG)transfer FADUM09)functions M "," Using \ref{fourierR}) ) and \ref{gwstatprop}) ), and taking into account the orthogonality property \ref{poltenorthog}) ), after straight forward calculations, we arrive at the following statistical expressions: = 0, )^2> = P_h(k) > = 0, )^2> = P_h(k) where we have introduced the transfer functions = _s | |^2."
+In the above expression dQ represents integration over the possible directions. of gravitational wave (Le. bdkd£Q)., In the above expression $d\Omega$ represents integration over the possible directions of gravitational wave (i.e. $d^3{\bmath{k}} = k^2dkd\Omega$ ).
+ From (12..13)) and (20)) it. follows. that the transfer functions JP47.«(K) do not depend on time variable fo," From \ref{trans1_1}, \ref{trans1_2}) ) and \ref{transferfunction}) ) it follows that the transfer functions $\tilde{R}_{+1,+2}^2(k)$ do not depend on time variable $t$."
+ That results from the stationarity of the gravitational wave field., That results from the stationarity of the gravitational wave field.
+ The expressions for the transfer function are calculated in the similar wav as in (Baskaranetal.2008).. but now they are slightly more complicated: These transfer functions behave like At and. A? respectively when f+0.," The expressions for the transfer function are calculated in the similar way as in \citep{bppp08}, but now they are slightly more complicated: These transfer functions behave like $k^4$ and $k^6$ respectively when $k\rightarrow0$."
+ The statistical propertics of stochastic gravitational wave field may be characterized by the density. parameter Έλα (Allen1996)., The statistical properties of stochastic gravitational wave field may be characterized by the density parameter $\Omega_{gw}$ \citep{Allen1997}.
+.. Qa. is related to the power spectrum ο): PU) where kg=2afafeπίνωο. and Ho is the current Llubbleparameter.," $\Omega_{gw}$ is related to the power spectrum $P_h(k)$: = ( )^2P_h(k) where $k_H= 2\pi f_H/c = 2\pi H_0/c$, and $H_0$ is the current Hubbleparameter."
+" Phe density parameter Q,.. is the current dav ratio of energy. density of gravitational waves (per unit logarithmic interval in &) to the critical density. of the Universe poy=9οHSSEC.", The density parameter $\Omega_{gw}$ is the current day ratio of energy density of gravitational waves (per unit logarithmic interval in $k$ ) to the critical density of the Universe $\rho_{crit} = 3c^2H_0^2/8\pi G$.
+" For numerical estimations. we set Hubble parameter 44,=sr75luuAlpe and assume a simple power law spectrum for the density parameter ο] Oh)=O,, s)p."," For numerical estimations, we set Hubble parameter $H_0 = 75~\frac{{\rm
+km}}{{\rm sec}}/{\rm Mpc}$ and assume a simple power law spectrum for the density parameter $\Omega_{gw}$: (k)= (k_o)."
+ This form of spectrum can be usec as a good approximation for a large varicty of models in. frequency range of our interest., This form of spectrum can be used as a good approximation for a large variety of models in frequency range of our interest.
+ The flat. scale invariant power spectrum (also known as lHlarrison-Zeldovich power spectrum) corresponds to mp=0.," The flat, scale invariant power spectrum (also known as Harrison-Zeldovich power spectrum) corresponds to $n_T=0$."
+ To obtain mean —square deviations. we should integrate (17)) and (19)) with GAY spectrum. (24)).," To obtain mean square deviations, we should integrate \ref{nudotsquaremean}) ) and \ref{nuddotsquaremean}) ) with GW spectrum \ref{powerlawspectrum}) )."
+" The limits of integration Ay, ancl Fus are determined. from. following considerations: the highes requeney that puts in the effect is defined by cut-olf scale coming from the pulsar timing technique ancl determine w total time span of pulsar observations used to obtain rotational parameters. the lowest [requeney of CW that can »* probed. with that method. comes from the limitations of οΠΡdistance: fum= and luas=2r/elin. (note that our maximal (requeney coincides with minima requeney of usual pulsar timing searches lor GWs: in [ac Auinsαποαν.ax 1. for the sake of simplicity we used à= 1)."," The limits of integration $k_{\mathrm{min}}$ and $k_{\mathrm{max}}$ are determined from following considerations: the highest frequency that puts in the effect is defined by cut-off scale coming from the pulsar timing technique and determined by total time span of pulsar observations used to obtain rotational parameters, the lowest frequency of GW that can be probed with that method comes from the limitations of pulsar-Earth; $k_{\mathrm{min}}\approx\frac{2\pi}{D}$ and $k_{\mathrm{max}}=2{\pi}/cT_{{\mathrm{obs}}}$ (note that our maximal frequency coincides with minimal frequency of usual pulsar timing searches for GWs; in fact $k_{\mathrm{max}}=2\alpha{\pi}/cT_{{\mathrm{obs}}},~\alpha\simeq
+1$ , for the sake of simplicity we used $\alpha=1$ )."
+ After substitutions we obtain the following equations: and )s = o, After substitutions we obtain the following equations: )^2> = = and )^2> = =
+ After substitutions we obtain the following equations: and )s = oa, After substitutions we obtain the following equations: )^2> = = and )^2> = =
+Since the AD relation is not linear (Fig. 3)).,"Since the AD relation is not linear (Fig. \ref{fig3}) ),"
+" it is very natural to consider that the Re. R,. parameters are not the same for all stellar populations."," it is very natural to consider that the $R_\Sigma$, $R_{\sigma_r}$ parameters are not the same for all stellar populations."
+ The change of the ellipsoid axis ratio cannot explain the bending of the AD relation., The change of the ellipsoid axis ratio cannot explain the bending of the AD relation.
+ In external galaxies the change with colour of the density scale length 15 only about (de Jong. 1996)).," In external galaxies the change with colour of the density scale length is only about (de Jong, \cite{dj96}) )."
+ However the youngest stars may have a smaller density scale length close to the ISM one. though it also depends on the SER radial dependence.," However the youngest stars may have a smaller density scale length close to the ISM one, though it also depends on the SFR radial dependence."
+ More certain. the kinematic seale length of youngest populations must be close to the ISM kinematics with a nearly null gradient.," More certain, the kinematic scale length of youngest populations must be close to the ISM kinematics with a nearly null gradient."
+ The different rate of dynamical heating with radius should change the kinematics to the observed one for old populations., The different rate of dynamical heating with radius should change the kinematics to the observed one for old populations.
+ For these reasons. we have tried a modeling with different kinematic scale lengths for the bluest stars (bins 3-10) and for the reddest stars (bins 11-20).," For these reasons, we have tried a modeling with different kinematic scale lengths for the bluest stars (bins 3-10) and for the reddest stars (bins 11-20)."
+ The likelihood is improved: we find out ¢..—5.2kms! and the density scale length 1.5kpc., The likelihood is improved; we find out $v_\odot= 5.2\kms$ and the density scale length $1.8\kpc$.
+" We find for the bluest stars a nearly null kinematic. gradient. in accordance to our expectation. and for reddest stars R,. Is 9.7kpe (or Rye=L9 kpo) close to the Lewis Freeman ""s value (1989)) Πο=LLkpe based on 600 giants towards the centre and anticentre directions."," We find for the bluest stars a nearly null kinematic gradient, in accordance to our expectation, and for reddest stars $R_{\sigma_r}$ is $9.7\kpc$ (or $R_{\sigma_r^2}=4.9\kpc$ ) close to the Lewis Freeman 's value \cite{lf89}) ) $R_{\sigma_r^2}=4.4\kpc$ based on 600 giants towards the centre and anticentre directions."
+ The comparison of the maximum likelihood (ML) of respective models allows us to determine their relative goodness of fit: the likelihood ratio test giving a measure of the reliability of adding new parameters., The comparison of the maximum likelihood (ML) of respective models allows us to determine their relative goodness of fit: the likelihood ratio test giving a measure of the reliability of adding new parameters.
+ The likelthood-ratio statistics (Nemec Nemec 1991.. Kendall Stuart 1967)) 1s A=- where £4 is the likelihood for models based on I4 components.," The likelihood-ratio statistics (Nemec Nemec \cite{nn91}, Kendall Stuart \cite{ks67}) ) is $\lambda =
+ \frac{{\cal{L}}_1}{{\cal{L}}_2}$ where ${\cal{L}}_1$ is the likelihood for models based on $K_1$ components."
+ Models with a larger number of components A» may give a better fit with A smaller than one., Models with a larger number of components $K_2$ may give a better fit with $\lambda$ smaller than one.
+ The statistical significance of A is obtained by comparing 21ogA to the upper percentile of a X? distribution with the number of degrees of freedom equal to the difference in the number of parameters between the two models. (, The statistical significance of $\lambda$ is obtained by comparing $-2\log\lambda$ to the upper percentile of a $\chi^2$ distribution with the number of degrees of freedom equal to the difference in the number of parameters between the two models. (
+However. these results only hold if the asymptotic normality and efficiency of the ML estimator are satisfied. Kendall Stuart. 1967)).,"However, these results only hold if the asymptotic normality and efficiency of the ML estimator are satisfied, Kendall Stuart, \cite{ks67}) )."
+ Likelihoods of the discussed models are summarized in Tab. 3.., Likelihoods of the discussed models are summarized in Tab. \ref{tab-rep1}.
+ Summarizing: 1) We find that Gaussian fits are significantly improved when the vertex is adjusted (A.~220 for ~20 more parameters)., Summarizing: 1) We find that Gaussian fits are significantly improved when the vertex is adjusted $\lambda \sim 220$ for $\sim 20$ more parameters).
+ 2) Comparing the Gaussian fit with à non null vertex to Eq., 2) Comparing the Gaussian fit with a non null vertex to Eq.
+ | modelings. we see that the Gaussian gives a significantly improved fit (and considering they have more free adjusted parameters).," 1 modelings, we see that the Gaussian gives a significantly improved fit (and considering they have more free adjusted parameters)."
+ 3) More interesting is that when Eq., 3) More interesting is that when Eq.
+ | modelings are compared to the more similar Gaussian fits with null vertex. the likelihood 1s very significantly improved considering that the number of free parameters ts lowered by about 65.," 1 modelings are compared to the more similar Gaussian fits with null vertex, the likelihood is very significantly improved considering that the number of free parameters is lowered by about 65."
+ 4) Finally comparing the two Eq., 4) Finally comparing the two Eq.
+ | models with one or two kinematic scale lengths. the change of likelthoods ts significant (as we add one supplementary parameter. the fit is improved) 2log A=5.2 and 30 for 18 and 19 bins. corresponding respectively to an improvement with a probability of 0.98 and 1.5x10 ," 1 models with one or two kinematic scale lengths, the change of likelihoods is significant (as we add one supplementary parameter, the fit is improved) $-2\log\lambda$ =5.2 and 30 for 18 and 19 bins, corresponding respectively to an improvement with a probability of 0.98 and $1.-5*10^{-8}$ )."
+ Errors given in the last column of Tab., Errors given in the last column of Tab.
+ 2. are deduced from the diagonal of the covariance matrix., \ref{tab2} are deduced from the diagonal of the covariance matrix.
+ Existing. correlations raise the errors., Existing correlations raise the errors.
+ A good indicator is the global correlation coefficient (Eadie et al. 1971)) a quantity that i$ a measure of the total amount of correlation between a variable and all the other variables (this coefficient being one if the variable Is alinear combination of the other variables).," A good indicator is the global correlation coefficient (Eadie et al, \cite{ed71}) ) a quantity that is a measure of the total amount of correlation between a variable and all the other variables (this coefficient being one if the variable is alinear combination of the other variables)."
+" Most correlations and global correlations are small except for the followingvariables: 1) a correlation and a global correlation about 0.8 between ο. and Rv: 2) A; has correlations of 0.3 and 0.5 with ο, and Ry and a global correlation of 0.8; 3) 6. has correlations of 0.2 and 0.3 with ο. and #,, and a global correlation of0.5.", Most correlations and global correlations are small except for the followingvariables: 1) a correlation and a global correlation about 0.8 between $v_\odot$ and $R_\Sigma$; 2) $R_{\sigma_r}$ has correlations of 0.3 and 0.5 with $v_\odot$ and $R_\Sigma $ and a global correlation of 0.8; 3) $\alpha$ has correlations of 0.2 and 0.3 with $v_\odot$ and $R_{\sigma_r}$ and a global correlation of 0.5.
+ This means that errors given in Tab., This means that errors given in Tab.
+ 2 must be multiplied by ~1.52 forthe variables e.5.Ry anda.," \ref{tab2} must be multiplied by $\sim 1.5-2$ for the variables $v_\odot, R_{\sigma_r}, R_\Sigma $ and $\alpha$."
+ The estimated density scale length is directly proportional to the adopted solar galactic radius My (here taken as 8.5kpc)., The estimated density scale length is directly proportional to the adopted solar galactic radius $R_0$ (here taken as $8.5 \kpc$ ).
+ It also depends on the adopted Vy in a less direct way., It also depends on the adopted $V_0$ in a less direct way.
+ We find that a decrease of Vi by 10 percent increases the estimate of Rs by about 10 percent while the slope of the rotation curve Is unchanged. a=0.20.," We find that a decrease of $V_0$ by 10 percent increases the estimate of $R_\Sigma$ by about 10 percent while the slope of the rotation curve is unchanged, $\alpha=-0.23$."
+ Ry and Vy are not accurately known but they are linked by the determination of the τω ratio (from the Oort's constants difference ο 2B). it results that ARs remains nearly unchanged by the Ry and Vi; uncertainties.," $R_0$ and $V_0$ are not accurately known but they are linked by the determination of the $R_0$ $V_0$ ratio (from the Oort's constants difference $A-B$ ), it results that $R_{\Sigma}$ remains nearly unchanged by the $R_0$ and $V_0$ uncertainties."
+ It is worth noting that we find a decreasing rotation curve in agreement with the new determination of the Oort’s constants (Feast Whitelock. 1997.. or Olling Merrifield. 1998) (4|B=(VidHR)= 2.1kmis'kpe !)," It is worth noting that we find a decreasing rotation curve in agreement with the new determination of the Oort's constants (Feast Whitelock, \cite{fw97}, or Olling Merrifield, 1998) $A+B=(dV/dR)_0=-2.4\kms\kpc^{-1}$ )."
+" Considering the constraints given by the rotation curve (Rohlfs et al. 1988). à decreasing rotation curve is obtained only with small Ry and V, (smaller than the [AU recommended values)."," Considering the constraints given by the rotation curve (Rohlfs et al, 1988), a decreasing rotation curve is obtained only with small $R_0$ and $V_0$ (smaller than the IAU recommended values)."
+ Olling Merrifield (1998)) find Ry=T.1kpe and Vy=1sl1lans |., Olling Merrifield \cite{om98}) ) find $R_0=7.1\kpc$ and $V_0=184\kms$ .
+ The corresponding o. the logarithmic slope of the rotation curve. is then 0.10 not in excellent agreement with our local determinations.," The corresponding $\alpha$ , the logarithmic slope of the rotation curve, is then $-0.10$ not in excellent agreement with our local determinations."
+ This difference could be partly explained by the most evident bias of our model. te. the assumption of a null vertex," This difference could be partly explained by the most evident bias of our model, i.e. the assumption of a null vertex"
+because otherwise significant svstematic time lags between the 5-rav and X-ray [loving behaviour would be expected. contrary to the observations (e.g.. Harmanetal.(2001))).,"because otherwise significant systematic time lags between the $\gamma$ -ray and X-ray flaring behaviour would be expected, contrary to the observations (e.g., \cite{hartman01}) )."
+ In the present paper. we describe a newly developed combined SSC + ERC jet radiation transfer code. accounting for time-dependent particle acceleration and injection. radiative cooling. and escape. coupled to the sell-consistent treatment of the relevant photon emission. absorption. aud escape processes.," In the present paper, we describe a newly developed combined SSC + ERC jet radiation transfer code, accounting for time-dependent particle acceleration and injection, radiative cooling, and escape, coupled to the self-consistent treatment of the relevant photon emission, absorption, and escape processes."
+ In 82. we give a brief description of (he underlving blazar jet model., In \ref{model} we give a brief description of the underlying blazar jet model.
+ The munerical procedure used in our code will be outlined in 33.., The numerical procedure used in our code will be outlined in \ref{numerics}.
+ We present results of a detailed parameter study. relevant for application to intermediate and low-Irequency peaked BL Lac objects. in 84..," We present results of a detailed parameter study, relevant for application to intermediate and low-frequency peaked BL Lac objects, in \ref{results}."
+ We summarize in 35.., We summarize in \ref{summary}.
+ The blazar model used [or this study. is a generic leptonic jet model., The blazar model used for this study is a generic leptonic jet model.
+" It is assumed that a population of ultrarelativistic. non-thermal electrons (aud. positrons) is injected at a generally time-dependent rate into a spherical emitting volume of co-moving radius J, (the “blob”)."," It is assumed that a population of ultrarelativistic, non-thermal electrons (and positrons) is injected at a generally time-dependent rate into a spherical emitting volume of co-moving radius $R_b$ (the “blob”)."
+ The injected pair population is specified through an injection power μμ) and the spectral characteristics of the injected non-thermal electron distribution., The injected pair population is specified through an injection power $L_{\rm inj} (t)$ and the spectral characteristics of the injected non-thermal electron distribution.
+" We assume that electrons are injected. with a singlepower-law distribution with low and high energv cutoffs 54 and 725. respectively. and a spectral index 4 so that the injection function Q,(5:/) ""s 4]. in the co-moving frame of the emitting region. is with where V7 is the blob volume in the co-moving frame."," We assume that electrons are injected with a singlepower-law distribution with low and high energy cutoffs $\gamma_1$ and $\gamma_2$, respectively, and a spectral index $q$ so that the injection function $Q_e (\gamma; t)$ $^{-3}$ $^{-1}$ ], in the co-moving frame of the emitting region, is with where $V'_b$ is the blob volume in the co-moving frame."
+ The jet is powered by accretion of material onto a supermassive central object. which is accompanied by (the formation of an accretion disk.," The jet is powered by accretion of material onto a supermassive central object, which is accompanied by the formation of an accretion disk."
+ For the purpose of this studs. we have represented the disk bv a standard Shakura-Sunvaev disk wilh a bolometric luminosity of Lp=10” eres +.," For the purpose of this study, we have represented the disk by a standard Shakura-Sunyaev disk with a bolometric luminosity of $L_D = 10^{45}$ ergs $^{-1}$."
+ The choice of this and several other standard parameters is motivated by a recent modeling study of the LBL W Comae (Bottcher.Mukherjee.&Reimer2002)., The choice of this and several other standard parameters is motivated by a recent modeling study of the LBL W Comae \citep{bmr02}.
+. The randomly oriented magnetic field P is determined by an equipartition parameter ej. which is the fraction of the magnetic fiekl energy density jy compared (o its value for equipartition," The randomly oriented magnetic field $B$ is determined by an equipartition parameter $\epsilon_B$ , which is the fraction of the magnetic field energy density $u_B$ compared to its value for equipartition"
+Iu the phenomenological approach to decoliereuce. one defines à Markov map for the evolution of the reduced density matrix of the system.,"In the phenomenological approach to decoherence, one defines a Markov map for the evolution of the reduced density matrix of the system."
+ This map is expected to be an approximation of the true dviauies of a system interacting with an environment., This map is expected to be an approximation of the true dynamics of a system interacting with an environment.
+ Indeed. it can often be derived via some simplifviuge assumptious: typically in fact. like i the CaldeiraLegeet model. the euviromnent cousists of a system with infinitely many deerees of freedom. whose exact dynamics is too complicate to he dealt with exactly.," Indeed, it can often be derived via some simplifying assumptions: typically in fact, like in the Caldeira–Legget model, the environment consists of a system with infinitely many degrees of freedom, whose exact dynamics is too complicate to be dealt with exactly."
+ This approach has proven to be very effective: one can list among its ost inportant successes the explanation of the puzzliug behavior of quantum Schroddinger cats [0|]. and the emereeuce of classical properties in quautuim mechanics |1].., This approach has proven to be very effective: one can list among its most important successes the explanation of the puzzling behavior of quantum Schröddinger cats \cite{haro1} and the emergence of classical properties in quantum mechanics \cite{joos}.
+ For hese reasons. a few vears ago. the idea was proposed that decoherence might also explain the problem of quautwim chaos. that is. it could restore dynamical chaos iu the quanti evolution. that was missing iu pure systems evolving via he Schróddinger equation |13.11.19]..," For these reasons, a few years ago, the idea was proposed that decoherence might also explain the problem of quantum chaos, that is, it could restore dynamical chaos in the quantum evolution, that was missing in pure systems evolving via the Schröddinger equation \cite{tomas,andrey1,zurpaz}."
+ Actually. this approach las a louger üstory: as carly as 1981 Guarneri [1| had shown that allowing a random erturbation im the dynamics of a kicked rotor washed away the plienomenon of «παπα localization of classical chaos and Ott [12]. had worked out he parameter dependence of the eusuiug diffusion coeffideut D.," Actually, this approach has a longer history: as early as 1984 Guarneri \cite{italo1} had shown that allowing a random perturbation in the dynamics of a kicked rotor washed away the phenomenon of quantum localization of classical chaos and Ott \cite{ott} had worked out the parameter dependence of the ensuing diffusion coefficient $D$ ."
+ Dittrich and Crahaim [13].. IKolovsky [L1]. aud later Sundaram ct al.," Dittrich and Graham \cite{tomas}, , Kolovsky \cite{andrey1} and later Sundaram et al."
+ [23] have shown thatthe, \cite{arje1} have shown thatthe
+for 5«10° <b3.10 oor6«1091. 1.6«104E...,"for $5 \times 10^9$ $\le$ $\le 1.3 \times
+10^{12}$ or$6 \times 10^8$ $\le$ $\le
+1.6 \times 10^{11}$."
+ For T3«1052 oor 21610H L.. This calibrationis based on that of Ienuicutt(1998) but with the I&roupa(2002) IME. which vields SFRs a factor of 0.66 of those assuming the Salpeter(1955). IME.," For $> 1.3 \times 10^{12}$ or $> 1.6 \times 10^{11}$ , This calibrationis based on that of \citet{Kennicutt98} but with the \citet{Kroupa02} IMF, which yields SFRs a factor of $0.66$ of those assuming the \citet{Salpeter55} IMF."
+ To stunmarize. the recipe to estimate aand SER from a given set of 21 pau fux ane) redshift nieasureients is following.," To summarize, the recipe to estimate and SFR from a given set of 24 $\micron$ flux and redshift measurements is following."
+ Aun IDL implementation of the steps below is available at ourwebsite. (, An IDL implementation of the steps below is available at our. (
+"1) Check that the galaxy of iuterest does not harbor luminous AGN that dominates its 21 sau emission using the iu-IR (e.g...Doulevetal.2012) or X-ray criteria (ee.M by requiring Ly|0.58.0keV]<1072 eres)."" (","1) Check that the galaxy of interest does not harbor luminous AGN that dominates its 24 $\micron$ emission using the mid-IR \citep[e.g.,][]{Donley12} or X-ray criteria (e.g., by requiring $_{\rm
+ X}[0.5-8.0~{\rm keV}] < 10^{42}$ erg/s). ("
+"2) TTo calculatealenl fettotal"" VIRTR buninosity.luminosity. interpolateinter for forttl (2)⋅⋅ and B/(:) frou Tablect 2. and caleulate : the olject’se redshift.not thenaffect the ΓΑΠ substitute these values in Equation 5 alone with the observed 21 pau flux iu 11Js.","2) To calculate total IR luminosity, interpolate for the coefficients $A'(z)$ and $B'(z)$ from Table \ref{table_azbz} and calculate the luminosity distance for the object's redshift, then substitute these values in Equation \ref{eq5} along with the observed 24 $\micron$ flux in mJy."
+ The resulting LCPIBμις. is mut: of , The resulting $L({\rm TIR})_{{\rm new}}$ is in unit of. (
+‘£⋅∎ caleulate SFR (in ME.//vr). calculate tli areof5 nu) im 7.. then substitute it in Equation 8. at suppression due huuinosity range.,"3) To calculate SFR(in /yr), first calculate the rest-frame in using Equation \ref{eq7}, , then substitute it in Equation \ref{eq8} at the appropriate luminosity range."
+ region and Iu Section ??.. we discuss the validity of using 21 jan observation at 0«τς2.8 to estimate aand SFR: in Section ??.. we discuss the implications of the indicator iu lisht of the test results from) Section ??..," In Section \ref{sec:discuss_validity}, we discuss the validity of using 24 $\micron$ observation at $0 < z < 2.8$ to estimate and SFR; in Section \ref{sec:discuss_sfmodes}, we discuss the implications of the indicator in light of the test results from Section \ref{sec:indicator_test}."
+" Iu addition. we will discuss the implications of the new indicator on the masxinuun typical ffor high-: ULIRGsin Section ??.. labolsec:discuss,.alidit"," In addition, we will discuss the implications of the new indicator on the maximum typical for $z$ ULIRGs in Section \ref{sec:discuss_LEdd}."
+yT hes uccessofthenewindicator. wh framewacclength sof2 lautobjnuetz= Oto2.s. Aadicatesthatphotoimcetrgdoimiénatedbgarometiceiissisamspleuoyt ," The success of the new indicator, which utilizes the luminosity at rest-frame wavelengths of 24 $\micron$ to 6 $\micron$ at $z = 0$ to 2.8, indicates that photometry dominated by aromatic emission is a surprisingly good tracer for."
+"dy ΡΕRil ΣΕ ΠΡ res ubtinter themetallicitydependenceandthepresence f AGN,", We discuss this result in terms of two major factors affecting the aromatic feature strength: the metallicity dependence and thepresence of AGN.
+ The correlation of the aromatic luuinosity to wwas studied by Rieopoulouetal.(1999).. Rousseletal.(2001). and Elbazctal.(2002) usingFSO: later by λαοἳal.(2005) with Spitser: and recently by Elbazctal.(2011) withH," The correlation of the aromatic luminosity to was studied by \citet{Rigopoulou99}, \citet{Roussel01} and \citet{Elbaz02} using; later by \citet{Wu05} with ; and recently by \citet{Elbaz11} with."
+erschel Although lee scatter in L(s jan) lis apparent over a huge rauge of metallicity (Calzcttietal.2007).. for high-iuctallicity svsteuis (CZ>1/3 Z.. which is equivalent to 12| log(O/TT) z8.2). Calzetti reports that the stellar-coutimmun-subtracted PAT emissiou shows a good correlation with the SFR.," Although large scatter in L(8 $\micron$ is apparent over a large range of metallicity \citep{Calzetti07}, for high-metallicity systems $Z > 1/3$ $Z_{\odot}$, which is equivalent to $12 + $ $/$ H) $\gtrsim 8.2$ ), \citet{Calzetti10b} reports that the stellar-continuum-subtracted PAH emission shows a good correlation with the SFR."
+ Furthermore. Eneelbrachtetal.(2008) and Suüthotal.(2007). find that the ratio of aromatic huuiuositv to ddoes not vary sieuificautlv at ietallicity 21/3Z..," Furthermore, \citet{Engelbracht08} and \citet{JDSmith07} find that the ratio of aromatic luminosity to does not vary significantly at metallicity $\gtrsim 1/3~Z_{\odot}$."
+ Galaxies at :~2 have on average 0.3 dex lower ietalliitv than local galaxies (Erbetal.2006)., Galaxies at $z \sim 2$ have on average $0.3$ dex lower metallicity than local galaxies \citep{Erb06}.
+. Therefore. galaxies more Massive than 3.Lo? ML. should be sufficiently metal rich (>1/3Z.) that the (PALL ealibrationiscalid," Therefore, galaxies more massive than $3 \times 10^{9}$ $_{\odot}$ should be sufficiently metal rich $> 1/3~Z_\odot$ ) that the $/$ PAH calibration is valid."
+ Tf weconsidertheratiog faromatic δα functiono finetallicityiithele ftpanclof Figure’ withinthera ης z.theseatteroftherelationshipbetweenaromaticlumiuosituandL(TI Hsinfactaboutü.ldeinthenon SeygfevtiLENERsaimple," If we consider the ratio of as a function of metallicity in the left panel of Figure \ref{eightmicron_lir} within the range of metallicity expected at $z$, the scatter of the relationship between aromatic luminosity and is in fact about $0.1$ dex in the non-Seyfert/LINER sample."
+Thissuggestsaromeaticluminositytobcagex aathigh zgiceuthatanef fortisimadetoceclude AGN fromthesempl Suuthetal., This suggests aromatic luminosity to be a good indicator for at $z$ given that an effort is made to exclude AGN from the sample (such as the AGN exclusion criteria employed in tests in Sections \ref{sec:indicator_testindiv} and \ref{sec:indicator_teststat}) ).
+(2007) found that ACN can significantly suppress the PAID euiission (seealso.Moorwood1986:Rocheetal.1991:Cenzel 1998).," \citet{JDSmith07} found that AGN can significantly suppress the PAH emission \citep[see
+ also,][]{Moorwood86, Roche91, Genzel98}."
+". The suppression of PAID huuinositv is shown iu Fieure 7 (left) using the total PATD huninosities∙∙∙ measured by SmithSU(2007) and a iietallicity measurement frou Moustakasal.(2.20101,"," The suppression of PAH luminosity is shown in Figure \ref{eightmicron_lir} (left) using the total PAH luminosities measured by \citet{JDSmith07}
+ and a metallicity measurement from \citet{Moustakas10}."
+ Moderate huuinosits N.inosifv ACNhowever. coefficients do eiuissiou of5 the eutireM host⋅⋅ ealaxy.," Moderate luminosity AGN, however, do not affect the PAH emission of the entire host galaxy."
+nA the huuinositvu distance⋅⋅ for⋅ .aquedPandekePORN wu↴∖⋁⊲−⇁⊽meet IRSa n Diamondcompar unefan‘a NENDEr (2019)CY UM. DUC ov galaxies ⋅⋅aud. fouud that while the PAIL features in nuclear ⋅↴spectra are deeclearly ussuppressed. the: features⋅↴ £3) Tointhe uormal streueth in the outer disk.," \citet{DiamondStanic10} use IRS to compare nuclear and non-nuclear spectra of nearby Seyfert galaxies and found that while the PAH features in the nuclear spectra are clearly suppressed, the features are of normal strength in the outer disk."
+ Therefore. with rest-frame L(21 using PAID emission to trace £(TIR ) hosting moderate luminosity ACN. the PATIL. unsiug Equation to ACN is likely limited to the unclear the appropriate PATIL cussion arising from the rest of the ealaxy should still provide a good tracer of the SER.," Therefore, with regard to using PAH emission to trace in galaxies hosting moderate luminosity AGN, the PAH suppression due to AGN is likely limited to the nuclear region and PAH emission arising from the rest of the galaxy should still provide a good tracer of the SFR."
+ We further test the validity of aromatic emission (6.8. 8 qun rest-frame observation) as a measure of LCTIB) ina of IIIE-doninated galaxies with moderate to lieh metallicity iu the right panel of Figure 7..," We further test the validity of aromatic emission (e.g., 8 $\micron$ rest-frame observation) as a measure of in a sample of HII-dominated galaxies with moderate to high metallicity in the right panel of Figure \ref{eightmicron_lir}."
+ Iu this Figure. the 8 gan observations are from Spifzer aud the cestimates are based on theZRAS all-sky survey using 12. 25. 60. aud LOO jun observatious (Sandersetal. 2003)..," In this Figure, the 8 $\micron$ observations are from and the estimates are based on the all-sky survey using 12, 25, 60, and 100 $\micron$ observations \citep{Sanders03}. ."
+ This saaple is chosen such that it mimics the populationof star-forming galaxies expected at hieh 2: lugh-inctallicity normal star-forming galaxies without AGN., This sample is chosen such that it mimics the populationof star-forming galaxies expected at high $z$: high-metallicity normal star-forming galaxies without AGN.
+ The scatter of the relationship between L(S8 ju) and im Figure 7 is 0.11 dex with the ratio //L(GS piu) ituAO aa, The scatter of the relationship between ${8~\micron}$ ) and in Figure \ref{eightmicron_lir} is $0.14$ dex with the ratio ${8~\micron}$ ) of $4.3 \pm 1.6$ .
+iPbatesh2011) reports a similar LUTIR)//Lés jn) of 925 for theirHerschel nat SPL without applviug corrections for the shape of the SED (seealso.Nordonctal. 2011).," \citet{Elbaz11} reports a similar ${8~\micron}$ ) of $4.9^{+2.9}_{-2.2}$ for their sample out to $z \sim 3$ (that is,a $\sigma$ spread of $\sim$0.23 dex) without applying corrections for the shape of the SED \citep[see also,][]{Nordon11}. ."
+. Iu the process of constructing our iudieator. we have assumed that all non-local," In the process of constructing our indicator, we have assumed that all non-local"
+also used the C270 ancl D270 receivers to observe frequencies around 241802 MITzZ. which include several transitions of methanol.,"also used the C270 and D270 receivers to observe frequencies around $241802\,$ MHz, which include several transitions of methanol."
+" Observations were made in wobbler-switching mode. using a wobbling periodof 0.51Iz and a beam throw of £130""."," Observations were made in wobbler-switching mode, using a wobbling periodof $0.5\,$ Hz and a beam throw of $\pm 130\arcsec$."
+ The telescope pointing and focus were checked periodically throughout the observing run using the automated routines and on Maus. Jupiter ancl Venus.," The telescope pointing and focus were checked periodically throughout the observing run using the automated routines and on Mars, Jupiter and Venus."
+ Ou each receiver. we used one 256 channel 1 MIIz flilterbank in parallel with the Vespa correlator svstem. set to a resolution of 320 kIIz and a bandwidth of 160 MIIz (for the formaldehyde bands) or 320 MIIz (for the methanol band).," On each receiver, we used one $256\,$ channel $1\,$ MHz filterbank in parallel with the Vespa correlator system, set to a resolution of $320\,$ kHz and a bandwidth of $160\,$ MHz (for the formaldehyde bands) or $320\,$ MHz (for the methanol band)."
+ For all observations. lines were placed in the lower sideband (LSB) with the upper sideband (USB) suppressed at the 95—974. level.," For all observations, lines were placed in the lower sideband (LSB) with the upper sideband (USB) suppressed at the $95-97\%$ level."
+ Because the USD suppression is imperfect. we also took spectra of the image sidebands and all spectra were observed with at least two different local oscillator settings. with central frequencies chosen to avoid prominent lines in the USB.," Because the USB suppression is imperfect, we also took spectra of the image sidebands and all spectra were observed with at least two different local oscillator settings, with central frequencies chosen to avoid prominent lines in the USB."
+ We are confident Chat our spectra contain no contaminating features [rom the USD., We are confident that our spectra contain no contaminating features from the USB.
+ Most initial data reduction was accomplished bv the automated routines developed by IRAM., Most initial data reduction was accomplished by the automated routines developed by IRAM.
+ Data [rom the IRAM 30m telescope is presented to observers calibrated ancl in units of antenna temperature. T. corrected. for all atmospheric losses.," Data from the IRAM 30m telescope is presented to observers calibrated and in units of antenna temperature, ${\rm T_A^*}$, corrected for all atmospheric losses."
+ To convert [rom units of TA to units of flux. we can simplydivide by the gain. which is 0.13INJ«.! in the 21m band and 0.10IXJv. tin the 1.3 mim band.," To convert from units of ${\rm T_A^*}$ to units of flux, we can simplydivide by the gain, which is $0.13\,{\rm K\, Jy^{-1}}$ in the $2\,$ mm band and $0.10\,{\rm K\, Jy^{-1}}$ in the $1.3\,$ mm band."
+ Further data reduction was accomplished using the Contimuun and Line Analvsis Single-Dish Software (CLASS) package., Further data reduction was accomplished using the Continuum and Line Analysis Single-Dish Software (CLASS) package.
+" We made two distinct ivpes of observations: 1) integration centered on a(2000)=947""57°4.0(2000)13°16” (the location of IRC+10216) and 2) offset mapping observations."," We made two distinct types of observations: 1) integration centered on $\alpha(2000)=9^h47^m57^s.4, 
+\delta(2000)=13^{\circ}16^{\prime}44^{\prime\prime}$ (the location of IRC+10216) and 2) offset mapping observations."
+ The goal of the first tvpe of observation was simply to detect formaldehyde or methanol. while the mapping observations were undertaken after we detected formaldehyde and were aimed al determining the spatial cistvibution of the formaldehyde emission.," The goal of the first type of observation was simply to detect formaldehyde or methanol, while the mapping observations were undertaken after we detected formaldehyde and were aimed at determining the spatial distribution of the formaldehyde emission."
+ liteeration times for each spectrum are listed in Table 2.., Integration times for each spectrum are listed in Table \ref{inttime}.
+ The analvsis of each set of observations proceeded cdiflerentlv. so we discuss the analvses separatelv.," The analysis of each set of observations proceeded differently, so we discuss the analyses separately."
+ For the central position. all observations al a particular [frequency were summed {ο produce a total spectrum for each baud. aud a linear baseline was subtracted.," For the central position, all observations at a particular frequency were summed to produce a total spectrum for each band, and a linear baseline was subtracted."
+ These reduced spectra are shown in Figure 1.., These reduced spectra are shown in Figure \ref{cspec}.
+ The spectra shown from the C150 and D150 receivers are from the 1 MIIz filterbank. while the 241.8 GlIz spectrum comes Irom the Vespa correlator svstenm and has been smoothed once using the CLASS command. in order (o improve our signal to noise ratio.," The spectra shown from the C150 and D150 receivers are from the $1\,$ MHz filterbank, while the $241.8\,$ GHz spectrum comes from the Vespa correlator system and has been smoothed once using the CLASS command, in order to improve our signal to noise ratio."
+ All displaved line fits were produced using the method of line [itting in CLASS., All displayed line fits were produced using the method of line fitting in CLASS.
+ Line fits produced with the routine have four free parameters: area.," Line fits produced with the routine have four free parameters: area,"
+"2011),, and furthermore, that the measured duration could also be reproduced by a large range of secondary sizes with an appropriate combination of orbital eccentricity and w.",", and furthermore, that the measured duration could also be reproduced by a large range of secondary sizes with an appropriate combination of orbital eccentricity and $\omega$."
+ Yet we find that no blends fainter than AKp—3.5 give tolerable fits to the light curve (see Figure 2))., Yet we find that no blends fainter than $\Delta K\!p = 3.5$ give tolerable fits to the light curve (see Figure \ref{fig:backstar}) ).
+" A visual understanding of the underlying reason for this may be seen in the bottom panel of Figure 3,, in which we show a blend model that one would naively expect should be able to match the observations, according to the crude recipe described above."," A visual understanding of the underlying reason for this may be seen in the bottom panel of Figure \ref{fig:fits}, in which we show a blend model that one would naively expect should be able to match the observations, according to the crude recipe described above."
+" This particular blend scenario is marked with a cross in Figure 2,, and corresponds to Aó—5 and Mz= 1.0Mo, resulting in a magnitude difference of AKp=5.6 for the EB relative to the target."," This particular blend scenario is marked with a cross in Figure \ref{fig:backstar}, , and corresponds to $\Delta\delta = 5$ and $M_2 = 1.0$ $M_{\sun}$, resulting in a magnitude difference of $\Delta K\!p = 5.6$ for the EB relative to the target."
+" While this model does yield a good match to the measured depth, and even the total duration, it doesn’t perform nearly as well in the ingress/egress phases, which are too long when compared against the observations."," While this model does yield a good match to the measured depth, and even the total duration, it doesn't perform nearly as well in the ingress/egress phases, which are too long when compared against the observations."
+" The quality of this fit relative to the best planet fit, which can also be seen in the figure, corresponds to a 10.1c difference, and therefore rrejects it."," The quality of this fit relative to the best planet fit, which can also be seen in the figure, corresponds to a $\sigma$ difference, and therefore rejects it."
+" Thus, the reason most blends of this class can be ruled out is ultimately the high precision of the light curves, which provides a very strong constraint on the shape of the transit light curve, and in particular on the size ratio between the secondary and tertiary, which sets the duration of the ingress and egress phases."," Thus, the reason most blends of this class can be ruled out is ultimately the high precision of the light curves, which provides a very strong constraint on the shape of the transit light curve, and in particular on the size ratio between the secondary and tertiary, which sets the duration of the ingress and egress phases."
+" There is a very broad range of blends consisting of a background or foreground star transited by a planet (as opposed to a star) that are found by tto give satisfactory fits to the data, as shown in Figure 4.."," There is a very broad range of blends consisting of a background or foreground star transited by a planet (as opposed to a star) that are found by to give satisfactory fits to the data, as shown in Figure \ref{fig:back_plan}."
+ These viable blends occupy the area below the 3c contour represented with a thick white line., These viable blends occupy the area below the $\sigma$ contour represented with a thick white line.
+" Secondary stars of all spectral types (masses) are permitted, in principle, although in practice other constraints described below eliminate a substantial fraction of them."," Secondary stars of all spectral types (masses) are permitted, in principle, although in practice other constraints described below eliminate a substantial fraction of them."
+ All of these blends involve secondary+tertiary pairs that are within 4 magnitudes of the target in the ppassband (diagonal dashed line in the figure)., All of these blends involve secondary+tertiary pairs that are within 4 magnitudes of the target in the passband (diagonal dashed line in the figure).
+ The tertiary sizes in these blends range from 0.42Rjyp to Ryup-, The tertiary sizes in these blends range from $R_{\rm Jup}$ to $R_{\rm Jup}$.
+" Our oobservations set a lower limit of about Mo for the secondary masses of these blends, as described earlier; scenarios involving redder stars would result in transits at ssignificantly deeper than we observe (i.e., deeper than the measured depth + 3c)."," Our observations set a lower limit of about $M_{\sun}$ for the secondary masses of these blends, as described earlier; scenarios involving redder stars would result in transits at significantly deeper than we observe (i.e., deeper than the measured depth + $\sigma$ )."
+ Thisexclusion region is indicated by the shaded area., Thisexclusion region is indicated by the shaded area.
+" Additionally, blends that"," Additionally, blends that"
+each star along the LIB is essentially. determined. by the mass of the surrounding envelope that survives the mass-loss episodes during the RGB phase.,each star along the HB is essentially determined by the mass of the surrounding envelope that survives the mass-loss episodes during the RGB phase.
+ From the observational point of view. the first physical parameter allecting the LB morphology. is the metallicity.," From the observational point of view, the first physical parameter affecting the HB morphology is the metallicity."
+ Metal rich clusters usually have red LBs. while the metal-poor ones tend to have LBs populated at. blucr colours.," Metal rich clusters usually have red HBs, while the metal-poor ones tend to have HBs populated at bluer colours."
+ However. at. least. one aceitional. or second. parameter is required to explain the observed. colour distributions of LIB stars in Galactic GCs.," However, at least one additional, or second, parameter is required to explain the observed colour distributions of HB stars in Galactic GCs."
+ This second parameter problem has received much attention during the last four decades (see for example Gratton et al., This second parameter problem has received much attention during the last four decades (see for example Gratton et al.
+ 2010 and references therein)., 2010 and references therein).
+ Age has been the most commonly discussed. second parameter (Lee ct al., Age has been the most commonly discussed second parameter (Lee et al.
+ 1994: Dotter et al., 1994; Dotter et al.
+ 2010: Gratton et al., 2010; Gratton et al.
+ 9010)., 2010).
+ Very. recent. work by Gratton et al. (, Very recent work by Gratton et al. (
+2010) ancl Dotter et al. (,2010) and Dotter et al. (
+2010) shows that age as a global second parameter can account for &eneral LIB parameters such as mean colour and length of the LLB.,2010) shows that age as a global second parameter can account for general HB parameters such as mean colour and length of the HB.
+ Llowever. HB morphologies show a diversity which cannot be fully characterized by mean properties.," However, HB morphologies show a diversity which cannot be fully characterized by mean properties."
+ Sets of clusters with similar metallicitios (M3/MI3/MSO. see Ferraro et al.," Sets of clusters with similar metallicities (M3/M13/M80, see Ferraro et al."
+ 1998: 47 Tuc/NGC 6388/NCGC 6441. Hich et al.," 1998; 47 Tuc/NGC 6388/NGC 6441, Rich et al."
+ 1997) have LIBs distinguished. by gaps and long blue tails., 1997) have HBs distinguished by gaps and long blue tails.
+ Features like these have prompted: discussion. of additional parameters like mass-Ioss ellicicney during the RGB (Rood et al., Features like these have prompted discussion of additional parameters like mass-loss efficiency during the RGB (Rood et al.
+ 1993: Ferraro et al., 1993; Ferraro et al.
+ 1998: Dotter 2008) or initial He. abundance cüllerences (CD'AXntona et al., 1998; Dotter 2008) or initial He abundance differences (D'Antona et al.
+ 2002)., 2002).
+ Phe situation is still more complicated. because deviations from the general behaviour arise due to internal second or third. parameters that may vary from cluster to cluster., The situation is still more complicated because deviations from the general behaviour arise due to internal second or third parameters that may vary from cluster to cluster.
+ In this context NGC2SOS is a particulary interesting case which certainly requires an accurate ancl detailed comparison between observations ancl state-of-the-art theoretical models., In this context NGC2808 is a particulary interesting case which certainly requires an accurate and detailed comparison between observations and state-of-the-art theoretical models.
+ As shown in Figs 2 and 3.. the HB o£ NCC 2808 is characterized by à well populated red clump and. on the blue side. a well populated and extended: blue sequence reaching down to mpsssiy~21.5. almost 2 magnitudes fainter than the AIS turn-olL (PO).," As shown in Figs \ref{cmd} and \ref{cmdfuv}, the HB of NGC 2808 is characterized by a well populated red clump and, on the blue side, a well populated and extended blue sequence reaching down to $m_{F555W}\sim 21.5$, almost 2 magnitudes fainter than the MS turn-off (TO)."
+ ὃν combining observations in optical and far-UV. bands we can sample LB stars according to their. temperatures and spectral type., By combining observations in optical and far-UV bands we can sample HB stars according to their temperatures and spectral type.
+ Ehe strategy. adopted. here is to. select stars along the LLB in the most appropriate CMD according to their temperature., The strategy adopted here is to select stars along the HB in the most appropriate CMD according to their temperature.
+ Hence. the red. clump stars have been selected. in the optical CALD (mipsssiw. msasonΠΕ). that is more sensitive to relatively cool stars. while the hottest LLB stars were selected in the far-UV C'MD. where they are the most luminous sources.," Hence, the red clump stars have been selected in the optical CMD $m_{F555W}$, $m_{F336W}-m_{F555W}$ ), that is more sensitive to relatively cool stars, while the hottest HB stars were selected in the far-UV CMD, where they are the most luminous sources."
+ These selection criteria allow also a more accurate distinction between pure HB and. post-HID. stars., These selection criteria allow also a more accurate distinction between pure HB and post-HB stars.
+ The intermediate-tempoerature LB population has been selected. in both CALDs. and. it was used a link between the cool and the hot part of the LIB. as well as to celine normalization criteria and the relative incompleteness of the two samples.," The intermediate-temperature HB population has been selected in both CMDs, and it was used a link between the cool and the hot part of the HB, as well as to define normalization criteria and the relative incompleteness of the two samples."
+ Two under-populated narrow regions (or “gaps”) are apparent in Figures 2 and 3., Two under-populated narrow regions (or “gaps”) are apparent in Figures 2 and 3.
+ These gaps were previously identified by. Dedin et al. (, These gaps were previously identified by Bedin et al. (
+2000: hereafter BOO) at V—18.6 and V—20 and the location in our CAID (at Ετσι=18.55 and missi=20.05) Is well consistent. with previous estimates.,2000; hereafter B00) at $V=18.6$ and $V=20$ and the location in our CMD (at $m_{F555W}=18.55$ and $m_{F555W}=20.05$ ) is well consistent with previous estimates.
+ The two gaps split the LIB in 4 subgroups (if. we also consider the standard subdivision at the cool anc hot. sides of the RR Lyrac instability strip). that were named RITB. EBTI. EBT? and EBT3 by DOO (see also Castellani ct al.," The two gaps split the HB in 4 subgroups (if we also consider the standard subdivision at the cool and hot sides of the RR Lyrae instability strip), that were named RHB, EBT1, EBT2 and EBT3 by B00 (see also Castellani et al."
+ 2006: lannicola et al., 2006; Iannicola et al.
+ 2009. hereafter. 100).," 2009, hereafter I09)."
+ For sake of homogeneity with previous work of our group. we will use the nomenclature introduced by Dalessandro ct al. (," For sake of homogeneity with previous work of our group, we will use the nomenclature introduced by Dalessandro et al. ("
+2008). and we will denote the four sub-populations as red LIB (RILB). blue HD (LLB). extreme (191112) and blue-hook (BLK). respectively.,"2008), and we will denote the four sub-populations as red HB (RHB), blue HB (HB), extreme HB (EHB) and blue-hook (BHk), respectively."
+ The physical properties of these 4 sub-groups will be discussed in the following section., The physical properties of these 4 sub-groups will be discussed in the following section.
+ Because of well known vignetting problems allecting the upper-right corner of all FIGOBW images (see Sect. 2)).," Because of well known vignetting problems affecting the upper-right corner of all F160BW images (see Sect. \ref{data}) ),"
+" we restrict our analysis to stars with distances r<75"" from Cas.", we restrict our analysis to stars with distances $r<75\arcsec$ from $C_{\rm grav}$.
+ Lhe selected populations are shown with different svmbols in Figs., The selected populations are shown with different symbols in Figs.
+ 2 and 3., 2 and 3.
+ The total IIB sample in our catalogue is 616 stars: 256 RIB. 243 121112. 66 EMB and 51 DII sources.," The total HB sample in our catalogue is 616 stars: 256 RHB, 243 BHB, 66 EHB and 51 BHk sources."
+ They. correspond. respectively. to (41d3). (39cxολ 11c2)% and (9c1X of the total.," They correspond, respectively, to $(41\pm3)\%$, $(39\pm3)\%$, $(11\pm2)\%$ and $(9\pm1)\%$ of the total."
+ These ratios are in agreement with those reported by 109 and are compatible with those quoted by BOO., These ratios are in agreement with those reported by I09 and are compatible with those quoted by B00.
+ We found 4 RR bvrae in common between the variable star list. published. by Samus et al. (, We found 4 RR Lyrae in common between the variable star list published by Samus et al. (
+2009) and Corwin et al. (,2009) and Corwin et al. (
+2004) ancl our DIID sample.,2004) and our BHB sample.
+ Given the small number. their inclusion. doesn't effect the comparison with theoretical models performed in the following sections. nor the results obtained in this work.," Given the small number, their inclusion doesn't effect the comparison with theoretical models performed in the following sections, nor the results obtained in this work."
+ To fully appreciate the oddity of NGC 2808's LLB. it is useful to compare it to a “normal” blue tail cluster. AlSO.," To fully appreciate the oddity of NGC 2808's HB, it is useful to compare it to a “normal” blue tail cluster, M80."
+ Of the clusters we've studied. in the far-UV. ALSO has the most extended. BP.," Of the clusters we've studied in the far-UV, M80 has the most extended BT."
+ Hs CALD along with that of, Its CMD along with that of
+similar formation redshift. anc between halos of similar mass similar concentration.,"similar formation redshift, and between halos of similar mass similar concentration."
+ In the following. formation redshift is defined as the redshift when a halo's main progenitor contains half its final mass. as used by 7 and GSWO05.," In the following, formation redshift is defined as the redshift when a halo's main progenitor contains half its final mass, as used by \cite{Gao2004} and GSW05."
+ We linearly interpolate halo growth between simulation outputs to increase the precision with which this redshift can be estimated., We linearly interpolate halo growth between simulation outputs to increase the precision with which this redshift can be estimated.
+ To estimate halo concentration we take the measured Voy and Vias For each halo and solve Eq., To estimate halo concentration we take the measured $V_{\rm vir}$ and $V_{\rm max}$ for each halo and solve Eq.
+ 5 in ?.., 5 in \cite{Navarro1996}.
+ In Fig., In Fig.
+ 3. we show the result of this exercise., \ref{fig3} we show the result of this exercise.
+ We plot relative bias as defined above (again estimated. from. an ensemble of 10 shulllecl catalogues) for absolute magnitude limited central galaxy subsaniples split. by colour., We plot relative bias as defined above (again estimated from an ensemble of 10 shuffled catalogues) for absolute magnitude limited central galaxy subsamples split by colour.
+ οσο were the subsamples with the most pronounced cllects in Fie., These were the subsamples with the most pronounced effects in Fig.
+ 2 and the results from that figure are repeated here as solid lines., \ref{fig2} and the results from that figure are repeated here as solid lines.
+ Ehe other lines show how the relative bias is reduced when shullling preserves halo formation redshift or 1alo concentration in addition to halo mass. thus how much of the assembly. bias in the original simulated catalogue can »e represented. using these additional halo properties.," The other lines show how the relative bias is reduced when shuffling preserves halo formation redshift or halo concentration in addition to halo mass, thus how much of the assembly bias in the original simulated catalogue can be represented using these additional halo properties."
+" Note hat results for the ""all colour” and “all galaxv catalogues oesented in Figure 2 show dependences on these additional »xwameters that are much weaker than the extreme cases shown here. typically less than a few percent."," Note that results for the “all colour” and “all galaxy” catalogues presented in Figure 2 show dependences on these additional parameters that are much weaker than the extreme cases shown here, typically less than a few percent."
+ We thus omit hem for clarity., We thus omit them for clarity.
+ Interestingly. Fig.," Interestingly, Fig."
+ 3. shows that neither. formation redshift nor concentration encodes sullicient. information o account for the assembly bias of the simulated. galaxy catalogue., \ref{fig3} shows that neither formation redshift nor concentration encodes sufficient information to account for the assembly bias of the simulated galaxy catalogue.
+" Of the two parameters. formation redshift is the most successful. accounting for about of the assembly ras for faint red central galaxies (at AM,5logsh=17 he relative bias is reduced from 1.37 to 1.22) but only a ew percent of the assembly bias for bright blue. central galaxies."," Of the two parameters, formation redshift is the most successful, accounting for about of the assembly bias for faint red central galaxies (at $M_{\rm b_J}\!-5\log h = -17$ the relative bias is reduced from 1.37 to 1.22) but only a few percent of the assembly bias for bright blue central galaxies."
+ Eniploving concentration as the second. parameter is only about half as effective in reducing the relative bias., Employing concentration as the second parameter is only about half as effective in reducing the relative bias.
+ Concentration dependences can account for only a small fraction of the measured assembly bias., Concentration dependences can account for only a small fraction of the measured assembly bias.
+ Clearly. although the galaxy content of our simulated halos depends: only on their mass ancl their assembly history. there is some aspect of the assembly history which is not encoded in halo concentration or formation redshift and which correlates with large-scale environment.," Clearly, although the galaxy content of our simulated halos depends only on their mass and their assembly history, there is some aspect of the assembly history which is not encoded in halo concentration or formation redshift and which correlates with large-scale environment."
+ Fig., Fig.
+ 3 demonstrates that halo concentration ancl halo formation redshift do not encode the same information about large-scale. clustering and that neither provides the information needed for precise modelling of the large-scale clustering of galaxies., \ref{fig3} demonstrates that halo concentration and halo formation redshift do not encode the same information about large-scale clustering and that neither provides the information needed for precise modelling of the large-scale clustering of galaxies.
+ Fie., Fig.
+ 4 explores the relation of central galaxy properties to formation redshift in more detail., \ref{fig4} explores the relation of central galaxy properties to formation redshift in more detail.
+ It shows the distribution of formation redshift for dark matter halos in two mass ranges. logMy214.0 and logM=10.911.1 (hAL.).," It shows the distribution of formation redshift for dark matter halos in two mass ranges, $\log
+M_{\rm vir}\!>\!14.0$ and $\log M_{\rm vir}\!=\!10.9-11.1$ $(h^{-1}
+M_\odot)$."
+ In cach range we compare the distribution for all halos (solid lines) with those for subpopulations which are extreme in their central galaxy properties. either luminosity or colour.," In each range we compare the distribution for all halos (solid lines) with those for subpopulations which are extreme in their central galaxy properties, either luminosity or colour."
+ Dashecl ancl dotted: curves in the left panel correspond. to the tails containing the brightest. and he faintest central galaxies respectively. while in the right xul they refer instead to the bluest and. reddest. central ealaxies.," Dashed and dotted curves in the left panel correspond to the tails containing the brightest and the faintest central galaxies respectively, while in the right panel they refer instead to the bluest and reddest central galaxies."
+ Fig., Fig.
+ 4 shows the well known result that low-mass halos jwe a broad cistribution of formation redshifts. centred at 2~1.7 and extending from 223 to 2: 0.5. while high-mass jos formed. more recently. with a distribution centred. at 20.7 and with tails extending from z70.0 to z—1.5 7)..," \ref{fig4} shows the well known result that low-mass halos have a broad distribution of formation redshifts, centred at $z\!\sim\!1.7$ and extending from $z\!>\!3$ to $z\!<\!0.5$ , while high-mass halos formed more recently, with a distribution centred at $z\!\sim\!0.7$ and with tails extending from $z\!\sim\!0.0$ to $z\!\sim\!1.5$ \citep[see][]{Lacey1993}."
+ For low-mass halos there is à weak but significant shift in formation redshift distribution between those with bright and those with faint central galaxies: faint. central galaxies live in halos that formed. svstematically earlier thanthose of their brighter counterparts. (, For low-mass halos there is a weak but significant shift in formation redshift distribution between those with bright and those with faint central galaxies: faint central galaxies live in halos that formed systematically earlier thanthose of their brighter counterparts. (
+In this mass bin the mean,In this mass bin the mean
+" In this section we use the information on the physical parameters of PMS stars of various masses and ages to study how star formation has in 3346 over the past ~30 MMyr, since this isproceeded the time span that we can effectively and accurately probe with our method."," In this section we use the information on the physical parameters of PMS stars of various masses and ages to study how star formation has proceeded in 346 over the past $\sim 30$ Myr, since this is the time span that we can effectively and accurately probe with our method."
+" As mentioned above, the distribution of our bona-fide PMS stars in the H-R diagram reveals a shortage of objects with ages around ~4—8 MMyr, suggesting a likely gap or lull in star formation at that time."," As mentioned above, the distribution of our bona-fide PMS stars in the H–R diagram reveals a shortage of objects with ages around $\sim
+4-8$ Myr, suggesting a likely gap or lull in star formation at that time."
+ The presence of two so clearly distinct groups of stars with Ha excess can only be interpreted as the result of distinct star formation episodes., The presence of two so clearly distinct groups of stars with $\alpha$ excess can only be interpreted as the result of distinct star formation episodes.
+ Hillenbrand et al. (, Hillenbrand et al. (
+2008) and Hillenbrand (2009) have argued that random luminosity spreads apparent in the H-R diagram of star and young clusters are often erroneously forminginterpreted as regionstrue luminosity spreads and taken as indicative of true age spreads.,2008) and Hillenbrand (2009) have argued that random luminosity spreads apparent in the H–R diagram of star forming regions and young clusters are often erroneously interpreted as true luminosity spreads and taken as indicative of true age spreads.
+" This is clearly not the case here, since no random spread in the luminosity of a single age population could produce such a distinctive bimodal distribution in the H-R diagram."," This is clearly not the case here, since no random spread in the luminosity of a single age population could produce such a distinctive bimodal distribution in the H–R diagram."
+" As regards the age distribution of PMS stars in 3346, a histogram is shown in reffig2, where PMS ages are binned using a constant logarithmic step (a factor of 2) that better reflects the relative age uncertainties stemming from the comparison of model isochrones with the actual data."," As regards the age distribution of PMS stars in 346, a histogram is shown in \\ref{fig2}, where PMS ages are binned using a constant logarithmic step (a factor of 2) that better reflects the relative age uncertainties stemming from the comparison of model isochrones with the actual data."
+" The solid line in reffig2 gives the number of stars inside each age bin as a function of time, whereas the dot-dashed line provides an apparent value of the star formation rate, in units of stars per derived the number of in each bin by Myr,the width of bythe bin."," The solid line in \\ref{fig2} gives the number of stars inside each age bin as a function of time, whereas the dot-dashed line provides an apparent value of the star formation rate, in units of stars per Myr, derived by dividing the number of objects in each bin by the width of the bin."
+" dividingNote that at the extremesobjects of the distribution it becomes more difficult to assign an age to the stars, thus the first and last bin are drawn with a dotted line to indicate a larger uncertainty."," Note that at the extremes of the distribution it becomes more difficult to assign an age to the stars, thus the first and last bin are drawn with a dotted line to indicate a larger uncertainty."
+ the number of detected PMS stars in the range 0.4—4.0 that at the time of the observations had Ha excess emission at the 4c level or above.," the number of detected PMS stars in the range $0.4 - 4.0$ that at the time of the observations had $\alpha$ excess emission at the $4\,\sigma$ level or above."
+" One limitation is caused by photometric incompleteness at low masses, which makes it more difficult to detect faint PMS stars in crowded environments."," One limitation is caused by photometric incompleteness at low masses, which makes it more difficult to detect faint PMS stars in crowded environments."
+ Another effect is the uncertainty on the fraction of PMS stars that at any given time show excess Ha emission., Another effect is the uncertainty on the fraction of PMS stars that at any given time show excess $\alpha$ emission.
+" For younger PMS stars (S8 MMyr), whose position in the H-R diagram is well separated from that of field MS stars, this fraction can be estimated from the ratio of stars with and without Ha excess in the same of the diagram."," For younger PMS stars $\lesssim 8$ Myr), whose position in the H–R diagram is well separated from that of field MS stars, this fraction can be estimated from the ratio of stars with and without $\alpha$ excess in the same region of the diagram."
+ The data show that at the time of the regionobservations this ratio was 0.28+0.04 (see also IID)., The data show that at the time of the observations this ratio was $0.28 \pm 0.04$ (see also II).
+" The thin solid line in reffig2 (green in the online version) shows the age distribution of all stars in the H-R diagram younger than MMyr, shifted vertically by —0.56ddex and is in excellent agreement with the histogram of bona-fide PMS stars."," The thin solid line in \\ref{fig2} (green in the online version) shows the age distribution of all stars in the H–R diagram younger than Myr, shifted vertically by $-0.56$ dex and is in excellent agreement with the histogram of bona-fide PMS stars."
+" However, it is presently not known how this ratio would change at older ages, and it is also expected to depend on the mass of the stars, so at this stage it is to set a lower limit to the true star formation rate."," However, it is presently not known how this ratio would change at older ages, and it is also expected to depend on the mass of the stars, so at this stage it is only possible to set a lower limit to the true star formation rate."
+" onlyNear the possiblepeak of the distribution, at ~0.4 MMyr, this limit corresponds to ~200 "," Near the peak of the distribution, at $\sim 0.4$ Myr, this limit corresponds to $\sim 200$ "
+BLTP 2009/07 Pom 1.5em Q.5¢m 2.5em Quantum ehromodsynamies predicts that the interaction between its fundamental constituents. quarks and eluons. ean lead to different states of strongly interacting matter. dependent on ils temperature ancl barvon densitv.,"BI-TP 2009/07 2cm 1.5cm 0.5cm 2.5cm Quantum chromodynamics predicts that the interaction between its fundamental constituents, quarks and gluons, can lead to different states of strongly interacting matter, dependent on its temperature and baryon density."
+ We first survey. the possible states of matter in QCD and discuss the transition from a color-confining hadronic phase to a plasma of deconlfined colored quarks aud gluons., We first survey the possible states of matter in QCD and discuss the transition from a color-confining hadronic phase to a plasma of deconfined colored quarks and gluons.
+ Next. we summarize (he results from non-perturbative studies of QCD at finite temperature and barvon density. and address the origin of deconfinement in the different regimes.," Next, we summarize the results from non-perturbative studies of QCD at finite temperature and baryon density, and address the origin of deconfinement in the different regimes."
+ Finally. we consider possible probes to test the basic features of bulk matter in QCD.," Finally, we consider possible probes to test the basic features of bulk matter in QCD."
+"In summary, we estimate a distance of 270+120 pc forCoRoT-2A,, based on interstellar absorption; the error was estimated from the standard deviation of the distances of the stars with Kj EWs in a +5 bband around our measured EW (cf.,","In summary, we estimate a distance of $270\pm 120$ pc for, based on interstellar absorption; the error was estimated from the standard deviation of the distances of the stars with $_4$ EWs in a $\pm 5$ band around our measured EW (cf.,"
+ Fig. 5))., Fig. \ref{fig:welsh}) ).
+" hhas a close neighbor,J19270636+0122577,, about 4” in southeast direction."," has a close neighbor, about $4\arcsec$ in southeast direction."
+ ? found that the color of this visual companion is consistent with a late-K or early-M type star located at the same distance asCoRoT-2A., \citet{Alonso2008} found that the color of this visual companion is consistent with a late-K or early-M type star located at the same distance as.
+". During five of our 24 UVES observations this nearby neighbor was placed inside the slit along withCoRoT-2A,, which we used to obtain a low SNR spectrum of the companion."," During five of our 24 UVES observations this nearby neighbor was placed inside the slit along with, which we used to obtain a low SNR spectrum of the companion."
+ The separation of the two objects allowed us to separately extract the spectra using the UVES pipeline., The separation of the two objects allowed us to separately extract the spectra using the UVES pipeline.
+" Because the companion is ~3.5 mag fainter than iin the visual band, the resulting spectrum has an SNR of no more than 10—20 depending on wavelength."," Because the companion is $\approx 3.5$ mag fainter than in the visual band, the resulting spectrum has an SNR of no more than $10-20$ depending on wavelength."
+ It is dominated by absorption lines from neutral and singly-ionized metals., It is dominated by absorption lines from neutral and singly-ionized metals.
+" In particular, we find strong absorption in andMgu, whereas theCau lines are comparably weak."," In particular, we find strong absorption in and, whereas the lines are comparably weak."
+" Furthermore, we find a relatively weak Ha line and distinct edges caused by titanium oxide (TiO) absorption."," Furthermore, we find a relatively weak $\alpha$ line and distinct edges caused by titanium oxide (TiO) absorption."
+" The TiO bands are a valuable indicator of stellar effective temperature if the metallicity is known, while they are less sensitive to surface gravity (?).."," The TiO bands are a valuable indicator of stellar effective temperature if the metallicity is known, while they are less sensitive to surface gravity \citep{MiloneBarbuy1994}."
+" To obtain an estimate of the effective temperature, we compared the coadded companion spectrum to synthetic spectra calculated with (?) using line lists containing the TiO and ZrO lines compiled by?hosted/plez.."," To obtain an estimate of the effective temperature, we compared the coadded companion spectrum to synthetic spectra calculated with \citep{GrayCorbally1994} using line lists containing the TiO and ZrO lines compiled by \cite{Plez1998}."
+" Assuming solar metallicity, we set up a Markov-Chain Monte-Carlo (MCMC) framework to find an estimate of the effective temperature and the associated error."," Assuming solar metallicity, we set up a Markov-Chain Monte-Carlo (MCMC) framework to find an estimate of the effective temperature and the associated error."
+ We used uniform priors on the stellar parameters and allowed for an additional normalization constant accounting for inadequacies during blaze correction and continuum normalization., We used uniform priors on the stellar parameters and allowed for an additional normalization constant accounting for inadequacies during blaze correction and continuum normalization.
+ We focused on the strongest TiO band with its bandhead at 7054 aand analyzed the three absorption edges individually., We focused on the strongest TiO band with its bandhead at $7054$ and analyzed the three absorption edges individually.
+ The resulting three We tried to use SME to fit synthetic spectra to several spectral lines known to be sensitive to the stellar parameters., The resulting three We tried to use SME to fit synthetic spectra to several spectral lines known to be sensitive to the stellar parameters.
+" Our efforts were, however, strongly hampered by the low SNR of the spectral data owing to the faintness of the companion."," Our efforts were, however, strongly hampered by the low SNR of the spectral data owing to the faintness of the companion."
+"From the analysis of the doublet at 5890A,, the lines at 6122, 6162, and 6439A,, and a set of 137 single Fe lines, we found the stellar spectrum to be best described by effective temperatures around 4000 K and a surface gravity between 4.6 and 4.9.","From the analysis of the doublet at 5890, the lines at $6122$, $6162$, and $6439$, and a set of 137 single Fe lines, we found the stellar spectrum to be best described by effective temperatures around 4000 K and a surface gravity between $4.6$ and $4.9$."
+" Unfortunately, the quality of the spectrum made an analysis of the Ha line impossible."," Unfortunately, the quality of the spectrum made an analysis of the $\alpha$ line impossible."
+ The set of lines was also used to quantify the effect of rotational line broadening., The set of lines was also used to quantify the effect of rotational line broadening.
+" Neglecting additional line broadening effects, we obtained an upper limit of 10 oon vsini at a confidence level of 90 These findings are compatible with the companion being a K9 star (?,p.388,Table15.7).."," Neglecting additional line broadening effects, we obtained an upper limit of 10 on $v\sin{i}$ at a confidence level of $90$ These findings are compatible with the companion being a K9 star \citep[][p.~388, Table~15.7]{AllenEd4}."
+" Our results are in line with those of ? and ?,, who found the photometric magnitudes measured in optical (Exodat), near-infrared (2MASS), and infrared (Spitzer)) filter bands to be consistent with a late-K or early-M type companion star."," Our results are in line with those of \citet{Alonso2008} and \citet{Gillon2010}, who found the photometric magnitudes measured in optical (Exodat), near-infrared (2MASS), and infrared ) filter bands to be consistent with a late-K or early-M type companion star."
+" To find the radial velocity of the visual companion, we cross-correlated our five companion spectra with a template spectrum corresponding to our best-fit stellar parameters."," To find the radial velocity of the visual companion, we cross-correlated our five companion spectra with a template spectrum corresponding to our best-fit stellar parameters."
+" The radial velocity was estimated independently in our five companion spectra, and corrected for the wavelength drift visible in the telluric lines."," The radial velocity was estimated independently in our five companion spectra, and corrected for the wavelength drift visible in the telluric lines."
+" The resulting average radial velocity amounts to 23.9+0.4kms!,, a value close to CoRoT-2A'ss radial velocity of 23.2450.010 s!,, which we determined accordingly."," The resulting average radial velocity amounts to $23.9\pm0.4$, a value close to s radial velocity of $23.245\pm0.010$ , which we determined accordingly."
+ We note that it is also independent of the details of the chosen spectral model., We note that it is also independent of the details of the chosen spectral model.
+" Given the apparent distance in the sky of 4"" and a distance of about 270 pc the projected distance between aand the companion amounts to about 1100 AU.", Given the apparent distance in the sky of $4\arcsec$ and a distance of about 270 pc the projected distance between and the companion amounts to about 1100 AU.
+" Because iis basically solar-like in mass, Kepler's third law yields a lower bound for the orbital period of 40000 a, which gives an orbital velocity of up to 0.9 kms!.. Thus, the radial velocity found for the visual companionof iis compatible with the hypothesis of a gravitationally bound companion."," Because is basically solar-like in mass, Kepler's third law yields a lower bound for the orbital period of $40\,000$ a, which gives an orbital velocity of up to 0.9 Thus, the radial velocity found for the visual companionof is compatible with the hypothesis of a gravitationally bound companion."
+star (133 M... helinm core: IW02).,star (133 $\msun$ helium core; HW02).
+ We use the implicit lvdvodvuamical code KEPLER (Weaver. Ziniuerman. Woosley 1978) to mocel the eutire evolution of the star and the resulting light curves.," We use the implicit hydrodynamical code KEPLER (Weaver, Zimmerman, Woosley 1978) to model the entire evolution of the star and the resulting light curves."
+ Since KEPLER ouly iuplenieuts erav diffusive radiation transport with approximate deposition of energv by eanuuna rays frou radioactive decay of CONi aud ο (Eastinan. Woosley. Weaver 1993). the light. curves obtained can only be followed as long as the SN is optically thick an there is a reasonably well-defined plotosphere.," Since KEPLER only implements gray diffusive radiation transport with approximate deposition of energy by gamma rays from radioactive decay of $^{56}$ Ni and $^{56}$ Co (Eastman, Woosley, Weaver 1993), the light curves obtained can only be followed as long as the SN is optically thick an there is a reasonably well-defined photosphere."
+ Also durius the very carlicst stages of the supernova (shock break out). KEPLER underestimates the color temperature by about a factor 2 (Dlinnikov 2003).," Also during the very earliest stages of the supernova (shock break out), KEPLER underestimates the color temperature by about a factor 2 (Blinnikov 2003)."
+ Finallv. we eiiplov. a musing procedure to simulate. in 1-D. the effects of mixing by Bavleighi-Tavlor instabilities after the superuova shock. which was calibrated using spectra of observed SN types.," Finally, we employ a mixing procedure to simulate, in 1-D, the effects of mixing by Rayleigh-Taylor instabilities after the supernova shock, which was calibrated using spectra of observed SN types."
+ The main effect of this mixing for the light curves is the spreading of some Ni into the euvelope., The main effect of this mixing for the light curves is the spreading of some $^{56}$ Ni into the envelope.
+ A supernova can be brieht either because it makes a lot of radioactive 7?Ni (as in Type Li supernovac) or because it has a large low density envelope aud large radius (as in bright Type ID supernovaec), A supernova can be bright either because it makes a lot of radioactive $^{56}$ Ni (as in Type Ia supernovae) or because it has a large low density envelope and large radius (as in bright Type II supernovae).
+ More radioactivity eives more energy at late times when the superuova has expanded οποιο] to become trausparcut (R — Lot? cin)., More radioactivity gives more energy at late times when the supernova has expanded enough to become transparent (R $\sim$ $^{15}$ cm).
+ A lavecr initial radius means that the star experiences less adiabatic cooling as it expands to that radius. resulting in a ligher luminosity at carly times.," A larger initial radius means that the star experiences less adiabatic cooling as it expands to that radius, resulting in a higher luminosity at early times."
+ Tere. the most inportant factors in determining the resulting helt curves are the mass of the progenitor star and the efficicucy of dredge-up of carbon from the core iuto the hydrogen envelope durius or at the end of central heli burning.," Here, the most important factors in determining the resulting light curves are the mass of the progenitor star and the efficiency of dredge-up of carbon from the core into the hydrogen envelope during or at the end of central helium burning."
+ The specifies of the plivsical process eucouutered here are unique to primordial stars., The specifics of the physical process encountered here are unique to primordial stars.
+ Lacking initial metals. they lave to produce the material for the CNO cvcle themsclyes. through the svuthesis of IC by the triple-alpha process.," Lacking initial metals, they have to produce the material for the CNO cycle themselves, through the synthesis of $^{12}$ C by the triple-alpha process."
+ Just enough 12€ is produced to initiate the CNO evele and bring it iuto equilibrium: a mass fraction of 10.? when central hydrogen burning starts. and a mass fraction ~10*[ during Lydrogen-shell burning.," Just enough $^{12}$ C is produced to initiate the CNO cycle and bring it into equilibrium: a mass fraction of $10^{-9}$ when central hydrogen burning starts, and a mass fraction $\sim 10^{-7}$ during hydrogen-shell burning."
+ At these low values the entropy in the hydrogen shell remains barely above that of the core. and the steep entropv eracicut at the upper edge of the helimm core that is typical for metaleuriched heliuni-buruiug stars is absent.," At these low values the entropy in the hydrogen shell remains barely above that of the core, and the steep entropy gradient at the upper edge of the helium core that is typical for metal-enriched helium-burning stars is absent."
+ The high masses of the stars also imply that a large fraction of the pressure is due to radiation., The high masses of the stars also imply that a large fraction of the pressure is due to radiation.
+ This is well known to facilitate convection., This is well known to facilitate convection.
+ For these reasons. he central convection zoue during heliuu buruiug cau ect close. nip af. or even penetrate the lydrogen-rich avers.," For these reasons, the central convection zone during helium burning can get close, nip at, or even penetrate the hydrogen-rich layers."
+ Once such iuixiug of higl-tempcrature hydroseu and carbon occurs. the two components burn violeutly. aud even without this rapid reaction. the hydrogen buruine iu he CNO evele increases proportionately to the additional carbon.," Once such mixing of high-temperature hydrogen and carbon occurs, the two components burn violently, and even without this rapid reaction, the hydrogen burning in the CNO cycle increases proportionately to the additional carbon."
+ Thus imixius of material from the heliuni-buruiug core. Which has a carbon abuudauce of order unity. is able o raise the energy ecucration rate in the hvdroseu-buruiug shell by orders of magnitude over its intrinsic value.," Thus mixing of material from the helium-burning core, which has a carbon abundance of order unity, is able to raise the energy generation rate in the hydrogen-burning shell by orders of magnitude over its intrinsic value."
+ This mdxiug has two major effects on the PPSN xogenitor: first. if inercases the opacity and energy eoncration iu the envelope. leading to a red-eiaut structure or the presuperuova star. m which the radius increases * over au order of magnitude.," This mixing has two major effects on the PPSN progenitor: first, it increases the opacity and energy generation in the envelope, leading to a red-giant structure for the presupernova star, in which the radius increases by over an order of magnitude."
+ Second. it decreases the uass of the Πο core. consequently leading to a smaller uass of Ni being svuthesized aud a snaller explosion eucrev.," Second, it decreases the mass of the He core, consequently leading to a smaller mass of $^{56}$ Ni being synthesized and a smaller explosion energy."
+ The former effect increases the huninosity of the supernova. especially at carly times. while the latter effect can weaken it.," The former effect increases the luminosity of the supernova, especially at early times, while the latter effect can weaken it."
+ Just how uch mixing occurs is uncertain from stellar modeling at the present. though it has been studied by several authors IHeger. Woosley Waters 2000: Alavigo 2001).," Just how much mixing occurs is uncertain from stellar modeling at the present, though it has been studied by several authors Heger, Woosley Waters 2000; Marigo 2001)."
+ In the present paper we account for these uncertainties by euploviug differcut values of convective overshooting whose presence is clear. but whose quantitative interaction with the burning is uukuown.," In the present paper we account for these uncertainties by employing different values of convective overshooting – whose presence is clear, but whose quantitative interaction with the burning is unknown."
+ A suite of represeutative models is chosen to address the expected range of presuperuova models - from blue Sperelant progenitors with little or no mixing to well-muned red hypergiauts. which cau have pre-SN radii of 20 AU or more.," A suite of representative models is chosen to address the expected range of presupernova models - from blue supergiant progenitors with little or no mixing to well-mixed red hypergiants, which can have pre-SN radii of 20 AU or more."
+ These models are suniuized in Table 1. iu which the names refer to the mass of the progenitor star (in units of AL.) aud the weak (NV). iuteiunediate (D. or strong (8) level of of couvective overshoot.," These models are summarized in Table 1, in which the names refer to the mass of the progenitor star (in units of $\msun$ ) and the weak (W), intermediate (I), or strong (S) level of of convective overshoot."
+" Note. that iu all cases the progenitor star is nonrotatiug. and all models employ the same procedure to mix 7""?Ni iuto the envelope after the supernova shock."," Note, that in all cases the progenitor star is nonrotating, and all models employ the same procedure to mix $^{56}$ Ni into the envelope after the supernova shock."
+ The KEPLER code can be used to compute approximate light curves and has been validated both against wich nore complex and realistic codes such as EDDINGTON and observations of a prototypical Type ILP supernova. SN 1969L (Weaver anc Woosley 1980: Eastinan et al 1991).," The KEPLER code can be used to compute approximate light curves and has been validated both against much more complex and realistic codes such as EDDINGTON and observations of a prototypical Type II-P supernova, SN 1969L (Weaver and Woosley 1980; Eastman et al 1994)."
+ Its deficiency. ds that it is a sinele, Its deficiency is that it is a single
+This last value is in the lower range of the values measured for RY Tau.,This last value is in the lower range of the values measured for RY Tau.
+ It corresponds to the value of a slowly rotating T Tauri star., It corresponds to the value of a slowly rotating T Tauri star.
+ From this point of view the example of RY Tau is an extreme one., From this point of view the example of RY Tau is an extreme one.
+ There is a large discrepancy for RY Tau between determinations of the spectroscopic projected rotational velocity and of the photometric rotational period., There is a large discrepancy for RY Tau between determinations of the spectroscopic projected rotational velocity and of the photometric rotational period.
+" In this particular case a value of 52-55 kms-! is inferred for Vsini (Petrovet al.,, 1999;; Moraet al.,, 2001; Agra-Amboageet al.,, 2009)), where i is the angle between the line of sight and the stellar rotational axis."," In this particular case a value of $52$ $55\kmps$ is inferred for $V\sin i$ \citeauthor{Petrovetal99}, , \citeyear{Petrovetal99}; \citeauthor{Moraetal01}, \citeyear{Moraetal01}; \citeauthor{AgraAmboageetal09}, \citeyear{AgraAmboageetal09}) ), where $i$ is the angle between the line of sight and the stellar rotational axis."
+ This value of Vsini is among the highest measured for CTTS., This value of $V\sin i$ is among the highest measured for CTTS.
+" This specific determination of the projected rotational velocity seems to correspond to a special case of obscuration by the disk rather than a real rotational velocity, however."," This specific determination of the projected rotational velocity seems to correspond to a special case of obscuration by the disk rather than a real rotational velocity, however."
+" Thus, this value may not correspond to the rotation of the star itself."," Thus, this value may not correspond to the rotation of the star itself."
+" Bouvier et al. (1993,, 1995))"," Bouvier et al. \citeyear{Bouvieretal93}, \citeyear{Bouvieretal95}) )"
+" give a photometric period of 24 days corresponding to V5,=5.1kms, which is a lower limit."," give a photometric period of 24 days corresponding to $V_{\varphi,o}= 5.1\kmps$, which is a lower limit."
+" Flux measurements give a period from 5 to 66 days, closer to typical values for TTS (Petrov al.,, 1999))."," Flux measurements give a period from $5$ to $66$ days, closer to typical values for TTS \citeauthor{Petrovetal99}, \citeyear{Petrovetal99}) )."
+" However, the variation is far from periodic and again one cannot exclude veiling from the disk."," However, the variation is far from periodic and again one cannot exclude veiling from the disk."
+" In other words, rotation measurements are extremely difficult to obtain for this object."," In other words, rotation measurements are extremely difficult to obtain for this object."
+" We have tried to construct solutions with rotational values close to the canonical value of 15kms!, though within a factor of 2 or so, assuming that more precise observational data are needed."," We have tried to construct solutions with rotational values close to the canonical value of $15\kmps$, though within a factor of $2$ or so, assuming that more precise observational data are needed."
+ The topology of the streamlines and magnetic fieldlines of the solution in the poloidal plane is displayed in Fig. 4.1.., The topology of the streamlines and magnetic fieldlines of the solution in the poloidal plane is displayed in Fig. \ref{Fig1}.
+ The stellar jet is surrounded by a diskwind and is self-collimated by its own magnetic field., The stellar jet is surrounded by a diskwind and is self-collimated by its own magnetic field.
+ It is produced in a fairly narrow coronal hole of half opening angle 15 degrees., It is produced in a fairly narrow coronal hole of half opening angle 15 degrees.
+ The hole is surrounded by a large dead zone extending up to 8 stellar radii., The hole is surrounded by a large dead zone extending up to 8 stellar radii.
+" This is comparable to the dipolar structure of 1.2 kG around BP Tau reconstructed from ESPaDOnS observations (Donatietal.,2008).", This is comparable to the dipolar structure of 1.2 kG around BP Tau reconstructed from ESPaDOnS observations \citep{Donatietal08}.
+. We also plot in Fig., We also plot in Fig.
+ 4.2 the outflow speed along the polar axis and the one along the last open streamline connected to the star., \ref{Fig2} the outflow speed along the polar axis and the one along the last open streamline connected to the star.
+ The asymptotic speed along the polar axis is which is higher by a factor of 2 than our initial guess., The asymptotic speed along the polar axis is which is higher by a factor of 2 than our initial guess.
+" 'This may be quite large, but considering the uncertainties involved, it can be an acceptable value."," This may be quite large, but considering the uncertainties involved, it can be an acceptable value."
+ In Fig., In Fig.
+" 4.2 we plot the polar density (a), pressure (b) and temperature (c)."," \ref{Fig3} we plot the polar density (a), pressure (b) and temperature (c)."
+" The pressure is defined up to a constant P,, see Eq. (A9))."," The pressure is defined up to a constant $P_o$, see Eq. \ref{pressure}) )."
+" By varying this constant, we obtain the various curves of Figs."," By varying this constant, we obtain the various curves of Figs."
+" 4.2bb and c. Adding different values of P, only affects the asymptotic part of the curves.", \ref{Fig3}b b and c. Adding different values of $P_o$ only affects the asymptotic part of the curves.
+ Note that the temperature in Fig., Note that the temperature in Fig.
+" 4.2cc has a maximum along the polar axis, reaching a value of one million degrees, which might be a rather high value."," \ref{Fig3}c c has a maximum along the polar axis, reaching a value of one million degrees, which might be a rather high value."
+" However, this effective temperature of the plasma is calculated from the equation of state of the gas using the total pressure plotted in Fig."," However, this effective temperature of the plasma is calculated from the equation of state of the gas using the total pressure plotted in Fig."
+" 4.2bb. This total pressure may include in addition to the kinetic pressure, ram pressure or Alfvénn waves, etc."," \ref{Fig3}b b. This total pressure may include in addition to the kinetic pressure, ram pressure or Alfvénn waves, etc."
+" As we have shown in Aibéoetal.(2007),, the effective temperature can be easily ten times higher than the kinetic one with a relative amplitude of the Alfvénn waves 6B/B less than unity."," As we have shown in \cite{Aibeoetal07}, the effective temperature can be easily ten times higher than the kinetic one with a relative amplitude of the Alfvénn waves $\delta B /B$ less than unity."
+ Note that GómezdeCastro&Verdugo(2007) inferred from UV lines high electronic temperatures associated with a wind close to 10°K consistent with this scenario and an effectivetemperature of one million degrees., Note that \cite{GomezVerdugo07} inferred from UV lines high electronic temperatures associated with a wind close to $10^5$ K consistent with this scenario and an effectivetemperature of one million degrees.
+candidates that are persisteuthy Dhuninous iu and also bright enough to beobserved with resolution ultraviolet spectrograpls with good sigual-to- (S/N).,candidates that are persistently luminous in X-rays and also bright enough to beobserved with high-resolution ultraviolet spectrographs with good signal-to-noise (S/N).
+ The black hole is believed to be uudergomeg accretion. frou its D-star companion via Roche lobe overflow with au orbital period of 1.7 davs (C83: Soria et al., The black hole is believed to be undergoing accretion from its B-star companion via Roche lobe overflow with an orbital period of 1.7 days (C83; Soria et al.
+ 2001)., 2001).
+ Spectroscopic observatious of the D star indicate a large radial velocity scumiamplitude. Aj=235+ddkinos (C3).," Spectroscopic observations of the B star indicate a large radial velocity semiamplitude, $K_B=235 \pm 11 \ km \ s^{-1}$ (C83)."
+ Although LMC Νο has beeu intensively studied in the X-ray baud. UV spectroscopic observations of this target are more sparse.," Although LMC X-3 has been intensively studied in the X-ray band, UV spectroscopic observations of this target are more sparse."
+ The first (FUSE) observation of Ελ N-3. was iade by Witchines et al. (, The first ) observation of LMC X-3 was made by Hutchings et al. (
+2003. hereafter II03) in November 2001.,"2003, hereafter H03) in November 2001."
+ During the 21 ks exposure. LAIC δν was in its xiehtest XN-rav phase.," During the 24 ks exposure, LMC X-3 was in its brightest X-ray phase."
+ These observatious revealed a woad emission feature in the vicinity of the AALO31.9.1037.6 doublet. which IIutehiugs et al.," These observations revealed a broad emission feature in the vicinity of the $\lambda \lambda$ 1031.9,1037.6 doublet, which Hutchings et al."
+ argue is due to the blend of the 1031.9 and 1037.6 lues (soe them Figure [)., argue is due to the blend of the 1031.9 and 1037.6 lines (see their Figure 4).
+ From this ciission. thev neasured a shift in the comission velocity between two binary phases of about 100-150 kau +.," From this emission, they measured a shift in the emission velocity between two binary phases of about 100-150 km $^{-1}$."
+ Asstunine the eenission arises in the mner parts of the black holes accretion disk. they couclude that the ΙΙΙ mass of the B star aud the black hole are 13 aud 15 A... respectively.," Assuming the emission arises in the inner parts of the black hole's accretion disk, they conclude that the minimum mass of the B star and the black hole are 13 and 15 $M_{\odot}$, respectively."
+" However. their measurement of the velocity variation of eenission with binary phase suffers frou, some uucertainties."," However, their measurement of the velocity variation of emission with binary phase suffers from some uncertainties."
+ First. the ephemeris they used to deteriuue the phase is from more than 20 vears ago.," First, the ephemeris they used to determine the phase is from more than 20 years ago."
+ Over such a lone time bascline. the errors contained iu he ephemeris could be accummilated. resulting in phases hat are off by zO.l at the lo level.," Over such a long time baseline, the errors contained in the ephemeris could be accumulated, resulting in phases that are off by $\approx$ 0.1 at the $1 \sigma$ level."
+ Another source of uucertaimty is the Dnited phase coverage of their observations., Another source of uncertainty is the limited phase coverage of their observations.
+ Due to the modest S/N of theirFUSE data. IIutchiugs et al.," Due to the modest S/N of their data, Hutchings et al."
+ sumuned theiFUSE spectra iuto only wo phase bins centered at phases = 0.53 aud 0.70. which oulv loosely coustrains Aj. the velocity scuuampltude of the black hole.," summed their spectra into only two phase bins centered at phases = 0.53 and 0.70, which only loosely constrains $K_{bh}$, the velocity semiamplitude of the black hole."
+ Overall. these limitations preclude αν fru conchisions about the nature and implications of the eecnission.," Overall, these limitations preclude any firm conclusions about the nature and implications of the emission."
+ Ad updated ephemeris aud expanded UV spectroscopy with better orbital coverage are ueeded to better constrain the implications of the UV. cussion lines aud he characteristics of this ΑΠΟ svstem., An updated ephemeris and expanded UV spectroscopy with better orbital coverage are needed to better constrain the implications of the UV emission lines and the characteristics of this XRB system.
+ Iu this paper. we revisit LAIC A-3 with a new ephemeris aud uew UV. observations.," In this paper, we revisit LMC X-3 with a new ephemeris and new UV observations."
+ We update the ephemeris of LAIC X5 wed on high-resolution optical spectroscopy recently obtained with the 6.510 AMaegelan-Clay telescope., We update the ephemeris of LMC X-3 based on high-resolution optical spectroscopy recently obtained with the 6.5m Magellan-Clay telescope.
+ We also report new ultraviolet observations of LAIC N-3 roni two hiel-vesolition instruments., We also report new ultraviolet observations of LMC X-3 from two high-resolution instruments.
+ First. we present eniporal monitoring of the ccluission usingFUSE data with a much longer time vascline that provides more than three quarters of the LMC Νο orbit.," First, we present temporal monitoring of the emission using data with a much longer time baseline that provides more than three quarters of the LMC X-3 orbit."
+ With the extendedFUSE data. we detect narrow Cluission from the binary system aud measure the various iu velocity aud intensity of the cussion as a function of orbital phase.," With the extended data, we detect narrow emission from the binary system and measure the variations in velocity and intensity of the emission as a function of orbital phase."
+ Second. we complement the aanalvsis with new observatious of the AAL2Z38.8.12 12.8 doublet obtained with the Cosnüc Origins Spectrograph (COS) ou board the(LEST).," Second, we complement the analysis with new observations of the $\lambda \lambda$ 1238.8,1242.8 doublet obtained with the Cosmic Origins Spectrograph (COS) on board the."
+ The eenission is detected at high sienificauce in the COS data and provides corroborating evidence of the velocity and intensity variations of the highlv ionized cussion., The emission is detected at high significance in the COS data and provides corroborating evidence of the velocity and intensity variations of the highly ionized emission.
+ During the tine of theFUSE observations iu 2001. LAIC N-3 was also observed quasi-simultancously in the Nav bandpass with aud theTE): see Wangetal.(2005) for full details.," During the time of the observations in 2004, LMC X-3 was also observed quasi-simultaneously in the X-ray bandpass with and the; see \citet{wang05} for full details."
+ While our euipliasis in this paper is ou the UV omission lues. we also briefly conuueut on the affiliated N-ray observations.," While our emphasis in this paper is on the UV emission lines, we also briefly comment on the affiliated X-ray observations."
+ With the combination of this information. we investieate the following questions: (1) Does the cCluission originate iu the accretion disk. the ionized stellar wind. the ilnunünated surface of the D star. or a hot spot (e.9.. from the spot where overflowing Roche lobe material inuipacts the accretion disk)? (," With the combination of this information, we investigate the following questions: (1) Does the emission originate in the accretion disk, the ionized stellar wind, the illuminated surface of the B star, or a hot spot (e.g., from the spot where overflowing Roche lobe material impacts the accretion disk)? ("
+2) Does the ddoublet cuiission arise in the same region as the eenission. 1.¢c.. are the aud ffeatures physically associated?,"2) Does the doublet emission arise in the same region as the emission, i.e., are the and features physically associated?"
+ Fitty three spectra of the optical counterpart (Warren Penfold 1975) of the N-ray binary svsteni were obtained on the niehts of 2005 January 2021. 2007 December 2021. aud 2008 February 27 9 2008 Mach lL using the Maecllan Diunori Ikvocera Echelle (NIKE) spectrograph (Bernsteinetal.2003) on the 6.5 τι Magellan-Clay telescope at Las Campanas Observatory.," Fifty three spectra of the optical counterpart (Warren Penfold 1975) of the X-ray binary system were obtained on the nights of 2005 January 20–24, 2007 December 20–21, and 2008 February 27 – 2008 March 1 using the Magellan Inamori Kyocera Echelle (MIKE) spectrograph \citep{bernstein03} on the 6.5 m Magellan-Clay telescope at Las Campanas Observatory."
+ The imstrinent was used iu the standard dualbbeani mode with a 170«570 slit and the 2« biuniug mode., The instrument was used in the standard dual-beam mode with a $1\farcs 0\times5\farcs 0$ slit and the $2\times 2$ binning mode.
+ With a few exceptions. the typical exposure time was 15800 s. We ouly focus ou spectra from the blue arm. which has a waveleneth coverage of 31305110 and resolving power of R=33.000.," With a few exceptions, the typical exposure time was 1800 s. We only focus on spectra from the blue arm, which has a wavelength coverage of 3430–5140 and resolving power of $R=33,000$."
+ The radial velocities were determined from these new spectra using the task within IRAF and then fitted to a circular orbit model., The radial velocities were determined from these new spectra using the task within IRAF and then fitted to a circular orbit model.
+" The typical uncertainties of the radial velocities are 3.5 kins ft,", The typical uncertainties of the radial velocities are 3.5 km $^{-1}$.
+ Following the initial observatious of IIutehiugsctal. (2003).. Ελ Νο was observed again withFUSE in April 2001. and during the latter cinupaignu quasi-simultaneous observations were recorded with andARATE (Waneetal.2005).," Following the initial observations of \citet{hutchings03}, , LMC X-3 was observed again with in April 2004, and during the latter campaign quasi-simultaneous observations were recorded with and \citep{wang05}."
+ The 2001FUSE observations provided a total exposure time of 86 ks (zls longer than the originalFUSE observations obtained in H03).," The 2004 observations provided a total exposure time of 86 ks $\approx 4
+\times$ longer than the original observations obtained in H03)."
+ Όλο N-3 is a variable source. and iu April 2001 the average UV flux of the target turned out to be zzLx fainter in the relevant spectral regions compared to the UV dus when Uutchinesotal.(2003). observed itin November 2001.," LMC X-3 is a variable source, and in April 2004 the average UV flux of the target turned out to be $\approx 4 \times$ fainter in the relevant spectral regions compared to the UV flux when \citet{hutchings03} observed it in November 2001."
+ Cousequcuthy. the S/N ratios are comparable in theFUSE data obtained in 2001 and 2004.," Consequently, the S/N ratios are comparable in the data obtained in 2001 and 2004."
+ However. the addition of the later data set siguificautlyv improves the orbital coverage of the ddata.," However, the addition of the later data set significantly improves the orbital coverage of the data."
+ Boththe 2001 aud 2004 FUSE data were reduced with CALFUSE. version 3.0 (Dixonotal. 2007)..," Boththe 2001 and 2004 data were reduced with CALFUSE, version 3.0 \citep{dixon07}. ."
+ It is, It is
+Fig 1 shows total intensity images of00.6... while Fie 2 shows an image of the entire. field.,"Fig 1 shows total intensity images of, while Fig \ref{fig_g309_primary_beam} shows an image of the entire field."
+ Properties. of sources of note are given in Table 3.., Properties of sources of note are given in Table \ref{tab_g309_sourcelist}.
+ Source 3 corresponds to the rregion (120900150.737 (2:?)..," Source 3 corresponds to the region G309.548–0.737 \cite{hcs79,ch87b}."
+" The remnant is comprised. of two morphological components: firstly two roughly circular ares of emission to the south-cast ancl north-west which we subsequently refer to as the ‘shell. and secondly two bright. sharply curved ares to the south-west and north-east (the ""ears)."," The remnant is comprised of two morphological components: firstly two roughly circular arcs of emission to the south-east and north-west which we subsequently refer to as the `shell', and secondly two bright, sharply curved arcs to the south-west and north-east (the `ears')."
+ The two parts of each. component are cliametrically opposed. with respect to the SNR’s geometric centre. anc the two components are oriented at. position angles perpendicular to cach other.," The two parts of each component are diametrically opposed with respect to the SNR's geometric centre, and the two components are oriented at position angles perpendicular to each other."
+ Although no connecting structure is apparent. the (wo arces of the shell can be construed to form a single circular ring.," Although no connecting structure is apparent, the two arcs of the shell can be construed to form a single circular ring."
+ Within the remnant. to the north-cast of centre. is à slightly. extended source. which we designate625235.," Within the remnant, to the north-east of centre, is a slightly extended source which we designate."
+.. At the available resolution (20 aresee see “Table 2)). this source has a cometary appearance. with a tail trailing out to the west.," At the available resolution $\sim$ 20 arcsec – see Table \ref{tab_g309_snr}) ), this source has a cometary appearance, with a tail trailing out to the west."
+ In Fig is shown an image of mmade using the observations in ‘Table 1.. but including the sixth ATCA antenna. 3 km west of the track upon which the other five antennae are stationed.," In Fig \ref{fig_g309_snr_central} is shown an image of made using the observations in Table \ref{tab_g309_observations}, but including the sixth ATCA antenna, 3 km west of the track upon which the other five antennae are stationed."
+ This increases the maximum baseline to 75000. m. corresponding to a resolution of 5 aresec. (," This increases the maximum baseline to $\sim$ 5000 m, corresponding to a resolution of 5 arcsec. ("
+Note that this gives a significant gap in the &0 coverage between 1500 m and 5000 m. and. so is not appropriate for imaging the entire remnant.),"Note that this gives a significant gap in the $u-v$ coverage between 1500 m and 5000 m, and so is not appropriate for imaging the entire remnant.)"
+ At this higher resolution. the source breaks into a double source to the cast. and a fainter extended source to the west.," At this higher resolution, the source breaks into a double source to the east, and a fainter extended source to the west."
+ To the north of the remnant is a narrow column of emission., To the north of the remnant is a narrow column of emission.
+ Extending northwards from the SNIUs north-eastern ear. this column bends around to the west and then to the north again. before opening up into a broader diffuse region forming a semi-circular arc.," Extending northwards from the SNR's north-eastern ear, this column bends around to the west and then to the north again, before opening up into a broader diffuse region forming a semi-circular arc."
+ The are is coincident with the Lla nebula ROW SO (7:?2).. while within it is a point source coincident with the position of the double star LID 119796 (?:7).," The arc is coincident with the $\alpha$ nebula RCW 80 \cite{rcw60,gbg+88}, while within it is a point source coincident with the position of the double star HD 119796 \cite{hsn71,hs85}."
+ Between aand the north-castern ear. is a collimated feature extending away from the SNIUs centre. which we dub the jet.," Between and the north-eastern ear, is a collimated feature extending away from the SNR's centre, which we dub the `jet'."
+ There is a clistinct break in the emission from the north-eastern ear where it intersects the jet., There is a distinct break in the emission from the north-eastern ear where it intersects the jet.
+ The jet and the break both lie along the svmmoetry axis defined by the two ears., The jet and the break both lie along the symmetry axis defined by the two ears.
+ The bright incar Component of the jet is 1755 long., The bright linear component of the jet is 5 long.
+ At its south-west end the jet abruptly faces. although a faint continuation can yo seen extending to the SNRs centre.," At its south-west end the jet abruptly fades, although a faint continuation can be seen extending to the SNR's centre."
+ Bevoned the break in he north-eastern ear is faint emission extending 3 arcmin xvond the SNA., Beyond the break in the north-eastern ear is faint emission extending 3 arcmin beyond the SNR.
+ The south-west ear also shows a break. although it is ess clistinet. and does not lic along the svmuetry axis.," The south-west ear also shows a break, although it is less distinct, and does not lie along the symmetry axis."
+ Paint emission which might correspond to a Less distinct and less collimated counterpart to the jet is seen just within this ear., Faint emission which might correspond to a less distinct and less collimated counterpart to the jet is seen just within this ear.
+ Derivecl parameters for the SNR are given in Table 2.. and were determined. by methods ceseribed in. Paper L.," Derived parameters for the SNR are given in Table \ref{tab_g309_snr}, and were determined by methods described in Paper \nocite{gmg98}."
+" ""Total [ux density measurements of aare shown in Fable ο we exclude single dish observations (?:?:2).. which are confused. by emission [from RCW SO to the north."," Total flux density measurements of are shown in Table \ref{tab_g309_snr} – we exclude single dish observations \cite{dtg69,ccg75,dshj95}, which are confused by emission from RCW 80 to the north."
+" We consequently compute a spectral index for SNICA090.2 oof a=0.58£0.09. where S,xο."," We consequently compute a spectral index for SNR of $\alpha = -0.53\pm 0.09$, where $S_\nu \propto \nu^\alpha$."
+ The spectral index calculated. here is somewhat steeper than previous results (a=0.37: Clark et al. 1975a)).," The spectral index calculated here is somewhat steeper than previous results $\alpha = -0.37$; Clark et al. \nocite{ccg75}) ),"
+ probably due to confusion with ROW 80., probably due to confusion with RCW 80.
+ We note that the largest spatial scale sampled in our image is 17 aremin. only slightlv larger than the remnant.," We note that the largest spatial scale sampled in our image is 17 arcmin, only slightly larger than the remnant."
+ Although our data include. additional. spacings corresponding to scales of 13 and 10 arcmin. one could argue that we are missing some of the SNIUSs flux density. and that its spectrum is consequently [latter than that we have just determined.," Although our data include additional spacings corresponding to scales of 13 and 10 arcmin, one could argue that we are missing some of the SNR's flux density, and that its spectrum is consequently flatter than that we have just determined."
+ However. the ATCA lux density is greater than that extrapolated from. lower frequency data. and we argue that little [lux is missing.," However, the ATCA flux density is greater than that extrapolated from lower frequency data, and we argue that little flux is missing."
+ Although625235... ROW SO and the column feature are of low surface brightness. we can put rough constraints on their spectra. as listed in Table 3..," Although, RCW 80 and the column feature are of low surface brightness, we can put rough constraints on their spectra, as listed in Table \ref{tab_g309_sourcelist}."
+ The lack of resolution in the MOST image of Whiteoak Creen (?) prevents a caleulation of the jets S43 MlIZ tux density and hence of its spectrum., The lack of resolution in the MOST image of Whiteoak Green \shortcite{wg96} prevents a calculation of the jet's 843 MHz flux density and hence of its spectrum.
+ Images of polarized emission were produced as in Paper I. the cllects of bandwidth: depolarization being minimised," Images of polarized emission were produced as in Paper I, the effects of bandwidth depolarization being minimised"
+the secondary components are even larger and would lead to ereat velocity separatious between the primary aud the secondary stars. which have never been found by spectral observations.,"the secondary components are even larger and would lead to great velocity separations between the primary and the secondary stars, which have never been found by spectral observations."
+ Fourier analysis is applied to investieate the intrinsic oscillation of the LSPVs over the raw data sets, Fourier analysis is applied to investigate the intrinsic oscillation of the LSPVs over the raw data sets.
+ Usiug the piriuueters (H4/A and q) obtained from the W-D code. the mean deusities of the LSPVs are deduced aud they are cousisteut with the red giant phase.," Using the parameters $R_{1}/A$ and $q$ ) obtained from the W-D code, the mean densities of the LSPVs are deduced and they are consistent with the red giant phase."
+" Then the pulsation constants are obtained by using the classical equation Q=Pralpi/p.2 to identify the pulsation mode,", Then the pulsation constants are obtained by using the classical equation $Q=P_{\rm{pul}}(\rho_{1}/\rho_{\sun})^{1/2}$ to identify the pulsation mode.
+ For star 77.7795.29. one pulsating frequency is detected ancl it is caused by the first overtone radial pulsation. aud for star 77.8031.19. the onulv pulsating⊓ frequency| is caused by the secoud overtone radial pulsation.," For star 77.7795.29, one pulsating frequency is detected and it is caused by the first overtone radial pulsation, and for star 77.8031.42, the only pulsating frequency is caused by the second overtone radial pulsation."
+ These agree with the theoretical value for red giants aud the conclusion of Woodctal.(1999)., These agree with the theoretical value for red giants and the conclusion of \citet{woo99}.
+. The stellar properties of LSPVs are derived by usine he information frou the pulsating binary uodel., The stellar properties of LSPVs are derived by using the information from the pulsating binary model.
+" We caleulate the bolometric maguitucde. ""numositv. radius and mass of the primary star. aud fud some of them conflict very seriously with the evolutionary properties of red eiauts."," We calculate the bolometric magnitude, luminosity, radius and mass of the primary star, and find some of them conflict very seriously with the evolutionary properties of red giants."
+ Iu articular. the masses for the star 77.7795.29 are TPSSOSI.12 ave both less than LAS... aud their uuimnosities are too Ligh for such low masses.," In particular, the masses for the star 77.7795.29 and 77.8031.42 are both less than $M_{\sun}$, and their luminosities are too high for such low masses."
+ It is difficult to imagine how such low mass stars couk rave such Lege huninositics., It is difficult to imagine how such low mass stars could have such large luminosities.
+" The situation gets even worse if we apply the mass ratio and the value of ""L4f/(L4| L5) to caleulate the mass iux uninositv of the secoudary star.", The situation gets even worse if we apply the mass ratio and the value of $L_{1}/(L_{1}+L_{2})$ ” to calculate the mass and luminosity of the secondary star.
+ Therefore. the radial velocities aud the masses computed from the pulsating binary model do no aeree with some observations aud facts about rec eqauts.," Therefore, the radial velocities and the masses computed from the pulsating binary model do not agree with some observations and facts about red giants."
+ We conclude that the model “pulsating binary has some ceficicucies in dealing with the observed properties of LSPVs aud that the binary lwpothesis for explaining the LSPs secs unreasonable., We conclude that the model “pulsating binary” has some deficiencies in dealing with the observed properties of LSPVs and that the binary hypothesis for explaining the LSPs seems unreasonable.
+ The authors are very grateful to Prof. Peter Wood των very coustructive sugecstionuseo acl helpful discussion., The authors are very grateful to Prof. Peter Wood for very constructive suggestions and helpful discussion.
+ They also thauk the auouviious referee for helpful suggestionsOO to improve the work. sienificauth., They also thank the anonymous referee for helpful suggestions to improve the work significantly.
+ This research is supported bv the National Natural Science Foundation of China (NSFC) through eraut 10778601 and the Ministry of Science and Techuologv of the People’s Republic of China through eraut 2007CD815106., This research is supported by the National Natural Science Foundation of China (NSFC) through grant 10778601 and the Ministry of Science and Technology of the People's Republic of China through grant 2007CB815406.
+age. that can be determuned from the N-vayv spectrin simultaueouslv with the clectrou temperature. from the ratio of the fluxes in different spectral lines from the same species.,"age, that can be determined from the X-ray spectrum simultaneously with the electron temperature, from the ratio of the fluxes in different spectral lines from the same species."
+ Rakowskietal(2003) studied shocks of the supernova remnant DEN L71 at both optical and N-rav waveleusths., \cite{Rakowski2003} studied shocks of the supernova remnant DEM L71 at both optical and X-ray wavelengths.
+ The remnant has shock velocities between 555 and 9801., The remnant has shock velocities between 555 and 980.
+ Correcting for temperature equilibration effects; Balowskietal.(2003) showed that decreases with increasing shock velocity.," Correcting for temperature equilibration effects, \cite{Rakowski2003} showed that $T_{\rm e}/T_{\rm p}$ decreases with increasing shock velocity."
+" T/T,From a theoretical poiut of view. if is uot well determined how inch clectrou-proton temperature equilibration would to expect behind a shock front."," From a theoretical point of view, it is not well determined how much electron-proton temperature equilibration would to expect behind a shock front."
+" Eimpirical studies of shocks with uceligible cosmic-ray acceleration. show that electron aud proton temperatures are equilibrated for slow shocks (e;< 400 +)), whereas the degree of thermal equilibration decreases for faster shocks (Rakowskietal.2003:Cchavamuanct2007:vanAdelsbergetal. 2008)."," Empirical studies of shocks with negligible cosmic-ray acceleration, show that electron and proton temperatures are equilibrated for slow shocks $v_{\rm s}<$ 400 ), whereas the degree of thermal equilibration decreases for faster shocks \citep{Rakowski2003,Ghavamian2007,Adelsberg}. ."
+. More specifically a relation of T./Tyxcy? has been suggested for shocks with ος>1400 (Ghavauuianetal.2007).," More specifically, a relation of $T_{\rm e}/T_{\rm p}\propto v_{\rm s}^{-2}$ has been suggested for shocks with $v_s > 400$ \citep{Ghavamian2007}."
+ This directly πράος that Tis nearlv constaut closely behind the shock frout (i.c.apartfromtemperatureequilibrationeffects;equa-tion2inChavanuanetal. 2002).. incdependeut of the shock velocity (equation 1)).," This directly implies that $T_{\rm e}$ is nearly constant closely behind the shock front \citep[i.e. apart from temperature equilibration effects, equation 2 in][]{Ghavamian2002}, independent of the shock velocity (equation )."
+ The origin of this relation is attributed by Cthavamianetal.(2007) to so-called lower hwbrid waves ahead of the shock. caused bv a (moderate) cosmic-ray precursor.," The origin of this relation is attributed by \citet{Ghavamian2007} to so-called lower hybrid waves ahead of the shock, caused by a (moderate) cosmic-ray precursor."
+ Note that vanAdelsbereetal.(2008). could not. reproduce TTxοὓς ?. using published data of He-spectia as well. but with different inodels for imterpreting Πα spectra.," Note that \cite{Adelsberg} could not reproduce $T_{\rm e}/T_{\rm p}\propto v_{\rm s}^{-2}$ , using published data of $\alpha$ -spectra as well, but with different models for interpreting $\alpha$ spectra."
+ Bykov&Uvarov(1999). and Bykov(2001) point out that for Mach uuubers AL<VINHh.=13 the electrous ahead of the shock have a tvpical thermal velocity exceediug the shock velocity. which may lead. through non-esonaut interactions with larec-wuplitude turbulent fluctuations iu the shock transition reeion. to collisionless clectrou-heating aud acceleration.," \citet{Bykov1999} and \cite{Bykov2004} point out that for Mach numbers $M < \sqrt{m_{\rm p}/m_{\rm e}} = 43$ the electrons ahead of the shock have a typical thermal velocity exceeding the shock velocity, which may lead, through non-resonant interactions with large-amplitude turbulent fluctuations in the shock transition region, to collisionless electron-heating and acceleration."
+ Tuterestinely. assunüng a typical sound speed of the zuubient medium of 10 kings. M=13 corresponds to a shock speed of ~LOO kis. For a cosmic-ray accelerating shock it is even more complex. as the incoming material is already: preheated by a cosmic-ray precursor. leading to a lower Mach uuuber at the main shock.," Interestingly, assuming a typical sound speed of the ambient medium of 10 km/s, $M=43$ corresponds to a shock speed of $\sim 400$ km/s. For a cosmic-ray accelerating shock it is even more complex, as the incoming material is already preheated by a cosmic-ray precursor, leading to a lower Mach number at the main shock."
+ This could mean that for cosmic-ray dominated shocks the electrou-proton temperatures could be equilibrated again., This could mean that for cosmic-ray dominated shocks the electron-proton temperatures could be equilibrated again.
+ Tere. we investigate the clectrou-proton temperature equilibration behind several shock frouts of RCW 386. using Nav (XAIN-Newtou/AIOS) aud optical data.," Here, we investigate the electron-proton temperature equilibration behind several shock fronts of RCW 86, using X-ray (XMM-Newton/MOS) and optical data."
+ The shocks of the supernova remnant ROW 86 are particularly well suited for this study. as the ionization age of the plasma iu parts of the remmant is low: ~«10?emὃς (ee.Vinketal.1997).. this value is low compared to the ionization age of most other reumautsL," The shocks of the supernova remnant RCW 86 are particularly well suited for this study, as the ionization age of the plasma in parts of the remnant is low: $\sim5\times10^9~{\rm cm}^{-3}{\rm s}$ \citep[e.g. ][]{Vink1997}, this value is low compared to the ionization age of most other remnants."
+o The low value for mot makes the measured electron telmperature a amore direct indication for the electrou temperature at the shock front., The low value for $n_{\rm e}t$ makes the measured electron temperature a more direct indication for the electron temperature at the shock front.
+ Additionally. the shocks of ROW 86 have a large varicty in speeds. as the progenitor exploded iu its own wind-blown bubble (e.g.B.Vinketal.1997) aud parts of the remnant have already hit the dense rini of this bubble.," Additionally, the shocks of RCW 86 have a large variety in speeds, as the progenitor exploded in its own wind-blown bubble \citep[e.g. ][]{Vink1997} and parts of the remnant have already hit the dense rim of this bubble."
+ As a result. the southwest (SW) rini las slowed dow: its N-ray oenission is mostly thermal aud the shock velocity is ~ 500 citepChavanian200L..," As a result, the southwest (SW) rim has slowed down; its X-ray emission is mostly thermal and the shock velocity is $\sim$ 500 \\citep{Ghavamian2001}."
+ The X-ray cunission of the rortheast (NE) shock is dominated by X-ray svuchrotron cCluission (Vinketal.2006).. and its shock velocity is neasured to be 6000 + 2800 ," The X-ray emission of the northeast (NE) shock is dominated by X-ray synchrotron emission \citep{Vink2006}, and its shock velocity is measured to be 6000 $\pm$ 2800 \\citep{Helder2009}."
+ILowever. its proton temperature is ow. which is an dudicatioun for οποια cosmic-ray acceleration.," However, its proton temperature is low, which is an indication for efficient cosmic-ray acceleration."
+ Furthermore. the mostly sub-solar abuncdanuces of the hermal N-rav enusson at the shock front of ROW 86 indicate enission from shocked aubicut media (rather tying the fitted parameters of the X-ray spectrum to the phyvsies at the shock frout.," Furthermore, the mostly sub-solar abundances of the thermal X-ray emission at the shock front of RCW 86 indicate emission from shocked ambient medium \citep[rather than shocked ejecta, e.g., ][]{Vink1997, Borkowski2001, Rho2002}, tying the fitted parameters of the X-ray spectrum to the physics at the shock front."
+ Iu this study. we use both published aud new optical observations.," In this study, we use both published and new optical observations."
+ The observation taken from the SW of the remnant was previously published im (μανάταιetal.(2001). and the observation from the NE riu was published iu Welderetal.(2009)., The observation taken from the SW of the remnant was previously published in \cite{Ghavamian2001} and the observation from the NE rim was published in \cite{Helder2009}.
+. The parameters of the north. CN) and east CE) observations are taken from Chavanuanetal.(2007) aud the northwest (NW) location from Chavanuuan(1999)., The parameters of the north (N) and east (E) observations are taken from \cite{Ghavamian2007} and the northwest (NW) location from \cite{GhavamianPhD}.
+. The exact locations of the N and E spectra have been kiudlv provided to us by Parviz Cliavaniau (private communication)., The exact locations of the N and E spectra have been kindly provided to us by Parviz Ghavamian (private communication).
+ We also use parameters of a spectrum published iu Long&Dlair (1990).. this spectrum is taken from the northern region as well.," We also use parameters of a spectrum published in \cite{Long}, this spectrum is taken from the northern region as well."
+ Table 2 sunuuarizes the line widths used in the present paper., Table 2 summarizes the line widths used in the present paper.
+ The optical image of Fig., The optical image of Fig.
+ 1...has been obtained with VLT/FORS?2 as pre-inage for the spectral observation. both in 2007 (programme ID 079.D-0735).," has been obtained with VLT/FORS2 as pre-image for the spectral observation, both in 2007 (programme ID 079.D-0735)."
+ The image is a combination of three exposures of 200s. through a narrowband Πα filter (H.Alpha*83)).," The image is a combination of three exposures of 200s, through a narrowband $\alpha$ filter )."
+ To correct for stellar light. three similar naages through a narrowhan¢l Πα filter. shifted 1500 ((H_Alpha/4500+61)) have been taken.," To correct for stellar light, three similar images through a narrowband $\alpha$ filter, shifted 4500 ) have been taken."
+ Figure shows the location of the slit. used to obtain the spectra.," Figure shows the location of the slit, used to obtain the spectra."
+ The spectra of the southeast (SE) part of the remnant were taken in a single 2733 s exposure. using a 1200R eri aud a sslit.," The spectra of the southeast (SE) part of the remnant were taken in a single 2733 s exposure, using a 1200R grism and a slit."
+ This setup resulted in a spectral resolution of 350d. which is sufficient for resolving the broad component. however. the narrow component of the line probably remains unresolved.," This setup resulted in a spectral resolution of 350, which is sufficient for resolving the broad component, however, the narrow component of the line probably remains unresolved."
+ The data were reduced with staudard data recuctiou steps. as described iu elderetal.(2009).. resulting in the spectra shown in Fie.," The data were reduced with standard data reduction steps, as described in \cite{Helder2009}, resulting in the spectra shown in Fig."
+2.. The slit crossed several filaments (Fig. 1))," The slit crossed several filaments (Fig. ),"
+. aud we here preseut spectra of the outer and inner filaments., and we here present spectra of the outer and inner filaments.
+ We focus ou the line width of the broadcomponeut. as this is directly related to the post-shock proton temperature.," We focus on the line width of the broadcomponent, as this is directly related to the post-shock proton temperature."
+ We fit the spectra with two Caussian lines. convolvedwith the resolution of the instrament set-up.," We fit the spectra with two Gaussian lines, convolvedwith the resolution of the instrument set-up."
+ Table lists the line widths obtained from the castern 4) and western (SES) spectrum in the slit (Fig. 1))., Table lists the line widths obtained from the eastern $_{\rm out}$ ) and western $_{\rm in}$ ) spectrum in the slit (Fig. ).
+ As the, As the
+"""hvsical/ parameters for. the radio source. components were computed assuming equipartition condition ancl using standard formulae. (Pacholezyk1010)...",Physical parameters for the radio source components were computed assuming equipartition condition and using standard formulae \citep{pacho70}.
+ We. considered yarticles with energy. between sin —100 and ομως 600.," We considered particles with energy between $\gamma_{\rm min} =$ 100 and $\gamma_{\rm
+ max}=$ 600."
+ The high energy. cut-oll corresponds to the value for which he break [frequency in the observer frame occurs at. 2.4 Gllz. as obtained from the best fit to the mocel (Section 3.2).," The high energy cut-off corresponds to the value for which the break frequency in the observer frame occurs at 2.4 GHz, as obtained from the best fit to the model (Section 3.2)."
+ Proton and electron energy. densities are assumed. to »e equal. and the spectral index is a=0.7 (see Section 3.2).," Proton and electron energy densities are assumed to be equal, and the spectral index is $\alpha = 0.7$ (see Section 3.2)."
+ We assume that the volume V of the emitting regions is a xolate spheroid: where @ and b are the spheroid major and minor axes. and © is the filling factor.," We assume that the volume $V$ of the emitting regions is a prolate spheroid: where $a$ and $b$ are the spheroid major and minor axes, and $\phi$ is the filling factor."
+ In the case of the entire racio source. we consider @ =140 mas. ὁ =25 mas. and ó=1.," In the case of the entire radio source, we consider $a=$ 140 mas, $b=$ 25 mas, and $\phi = 1$."
+ With these assumptions. we obtain an equipartition magnetic field Jf.= 4 mCG. and a minimum energy. density Main=14-10° erg/em?.," With these assumptions, we obtain an equipartition magnetic field $H_{\rm eq} =$ 4 mG, and a minimum energy density $u_{\rm min} = 1.4 \times 10^{-6}$ $^{3}$."
+ In the case of component N and S we assume @=23 mas. b —15 mas and O=1. and we infer a magnetic licld of 7 mG. and a minimum energy. density of 25. ' °erg/cm'. in agreement with the values derived for the entire source In general the shape of the svnchrotron spectrum of active racio sources is the result. of the interplay between freshly injected. relativistic particles ancl energy. losses that cause a depletion of high-energv particles which results into a steepening of the spectrum at. high. frequencies.," In the case of component N and S we assume $a=23$ mas, $b=$ 15 mas and $\phi$ =1, and we infer a magnetic field of 7 mG, and a minimum energy density of $=$ $\times$ $^{-6}$ $^{3}$, in agreement with the values derived for the entire source In general the shape of the synchrotron spectrum of active radio sources is the result of the interplay between freshly injected relativistic particles and energy losses that cause a depletion of high-energy particles which results into a steepening of the spectrum at high frequencies."
+ The discovery of a GPS radio source. considered to be à voung object where the radio emüssion started. a few. thousand vears ago. but displaving a steep spectral index (a> 1.0) in all its components is somewhat. surprising.," The discovery of a GPS radio source, considered to be a young object where the radio emission started a few thousand years ago, but displaying a steep spectral index $\alpha_{\rm 1.6}^{8.4} > 1.0$ ) in all its components is somewhat surprising."
+ Indeed. the strong steepening found across the components of LI5IS|O47 is incompatible with the moderate hieh- steepening predicted by a continuous injection of particles. suggesting that no injection/acceleration of fresh relativistic particles is currently occurring in any region," Indeed, the strong steepening found across the components of 1518+047 is incompatible with the moderate high-frequency steepening predicted by a continuous injection of particles, suggesting that no injection/acceleration of fresh relativistic particles is currently occurring in any region"
+"and at least as good as the well-known improved stabilized Conjugate Gradient (BiCGstab) method used,e.g., in ?..","and at least as good as the well-known improved stabilized Bi-Conjugate Gradient (BiCGstab) method used,e.g., in \citet{Yorke:2002p1}."
+" The general idea of minimal residual solvers based on the KSP method is the following: The i"" KSP is defined as Κι=spanib,Ab,A?b,waACT p)."," The general idea of minimal residual solvers based on the KSP method is the following: The $i^\mathrm{th}$ KSP is defined as $K_i = span\{\vec{b}, A \vec{b}, A^2 \vec{b}, ..., A^{i-1} \vec{b}\}$ ."
+ In each solver iteration i the GMRES method increments the subspace K; used with an additional basis vector A’! and approximates the solution of the system of linear equations by the vector X; which minimizes the norm of the residual r=|AX;—pl., In each solver iteration $i$ the GMRES method increments the subspace $K_i$ used with an additional basis vector $A^{i-1} \vec{b}$ and approximates the solution of the system of linear equations by the vector $\vec{x}_i$ which minimizes the norm of the residual $r = |A \vec{x}_i - \vec{b}|$.
+ This method converges monotonically and theoretically reaches the exact solution after performing as many iterations as the column number of the matrix A (which equals the number of grid cells)., This method converges monotonically and theoretically reaches the exact solution after performing as many iterations as the column number of the matrix $A$ (which equals the number of grid cells).
+" Of course, in practice the iteration is already stopped after reaching a specified relative or absolute tolerance of the residual, which normally takes only a small number of iterations."," Of course, in practice the iteration is already stopped after reaching a specified relative or absolute tolerance of the residual, which normally takes only a small number of iterations."
+ The computation of each iteration grows like O(P)., The computation of each iteration grows like $O\left(i^2\right)$.
+" In the current implementation, we use the so-called 'GMRES restarted' by default."," In the current implementation, we use the so-called 'GMRES restarted' by default."
+ GMRES restarted never performs all the iterations to reach the exact solution., GMRES restarted never performs all the iterations to reach the exact solution.
+" After a priorily fixed number of internal iterations n, the solver starts a second time in the first subspace Κι but with the last approximate solution χι."," After a priorily fixed number of internal iterations $n$, the solver starts a second time in the first subspace $K_1$ but with the last approximate solution $\vec{x_n}$."
+ Due to the growing computational effort with o(P) this approach generally results in a speedup of the computation., Due to the growing computational effort with $O\left(i^2\right)$ this approach generally results in a speedup of the computation.
+" The radiation transport module is parallelized for distributed memory machines, using the message passing interface (MPI) language."," The radiation transport module is parallelized for distributed memory machines, using the message passing interface (MPI) language."
+" The results of a detailed parallel performance test of the whole radiation transport module, including this GMRES solver is presented in Sect. 4.."," The results of a detailed parallel performance test of the whole radiation transport module, including this GMRES solver is presented in Sect. \ref{sect:performance}."
+ The approximate radiation transport introduced in the previous section can now be tested for realistic dust opacities in a standard benchmark test for irradiated circumstellar disk models., The approximate radiation transport introduced in the previous section can now be tested for realistic dust opacities in a standard benchmark test for irradiated circumstellar disk models.
+" The setup of the following comparison is adopted from ? and includes a central solar-type star, an irradiated circumstellar flared disk and an envelope."," The setup of the following comparison is adopted from \citet{Pascucci:2004p39}
+ and includes a central solar-type star, an irradiated circumstellar flared disk and an envelope."
+" We have to choose a low-mass central star, because no benchmark for high-mass stars was performed so far."," We have to choose a low-mass central star, because no benchmark for high-mass stars was performed so far."
+" However, the tests should not depend on the actual size and luminosity of the central star."," However, the tests should not depend on the actual size and luminosity of the central star."
+" For comparison, we choose a standard full frequency dependent Monte-Carlo based radiation transport code."," For comparison, we choose a standard full frequency dependent Monte-Carlo based radiation transport code."
+ The comparison is done for two different (low and high) optical depths taken from the original radiation benchmark test., The comparison is done for two different (low and high) optical depths taken from the original radiation benchmark test.
+" To test each of the components (gray and frequency dependent irradiation as well as Flux Limited Diffusion) of the proposed approximate radiative transfer separately, we perform several test runs with and without the different physical processes (absorption and re-emission) with both the Monte-Carlo based (see table 1)) and the approximate (see table 2)) radiative transfer code."," To test each of the components (gray and frequency dependent irradiation as well as Flux Limited Diffusion) of the proposed approximate radiative transfer separately, we perform several test runs with and without the different physical processes (absorption and re-emission) with both the Monte-Carlo based (see table \ref{MC-runs}) ) and the approximate (see table \ref{run-table}) ) radiative transfer code."
+" The stellar parameters are solar-like: The effective temperature of the star is 5800 K (T.=5800K) and the stellar radius is fixed to 1 solar radius (R,=1Ro).", The stellar parameters are solar-like: The effective temperature of the star is 5800 K $\left(T_* = 5800 \mbox{ K}\right)$ and the stellar radius is fixed to 1 solar radius $\left(R_* = 1 \mbox{ R}_\odot\right)$.
+ The disk ranges from rgig=1 AU up to rmax=1000 AU., The disk ranges from $r_\mathrm{min} = 1$ AU up to $r_\mathrm{max} = 1000$ AU.
+" Although the numerical setup of the gas density is done in spherical coordinates, the analytic setup of the gas density, as described in the original benchmark test, is given in cylindrical coordinates: with the radially and vertically dependent functions and making use of the following abbreviations The lowest density is limited to a relative factor of 10719? compared to the highest density (at rmin in the midplane) to avoid divisions by zero (e.g. in the calculation of the diffusion coefficients)."," Although the numerical setup of the gas density is done in spherical coordinates, the analytic setup of the gas density, as described in the original benchmark test, is given in cylindrical coordinates: with the radially and vertically dependent functions and making use of the following abbreviations The lowest density is limited to a relative factor of $10^{-100}$ compared to the highest density (at $r_\mathrm{min}$ in the midplane) to avoid divisions by zero (e.g. in the calculation of the diffusion coefficients)."
+ The normalization oo of the density setup is chosen to define different optical depths 7559044 through the midplane of the corresponding circumstellar disk (ata visual wavelength of 550 nm): The opacity tables used are the same as in the original benchmark test (?) taken from ?..," The normalization $\rho_0$ of the density setup is chosen to define different optical depths $\tau_{550\mathrm{nm}}$ through the midplane of the corresponding circumstellar disk (ata visual wavelength of $550$ nm): The opacity tables used are the same as in the original benchmark test \citep{Pascucci:2004p39}
+ taken from \citet{Draine:1984p594}. ."
+ They are displayed in Fig. 2.., They are displayed in Fig. \ref{Opacities}. .
+"Overall, the flow pattern is one in which, in each half wavelength in longitude, there is a local ’roll’ with E-W axis, the sense of which reverses in the next half wavelength.","Overall, the flow pattern is one in which, in each half wavelength in longitude, there is a local 'roll' with E-W axis, the sense of which reverses in the next half wavelength."
+" Within each roll, flow toward the equator acquires a negative component relative to the rotating frame, flow toward the pole, a positive component."," Within each roll, flow toward the equator acquires a negative component relative to the rotating frame, flow toward the pole, a positive component."
+" This means that locally the circulation in the roll is tending to conserve angular momentum, only weakly opposed by any longitudinal fluid pressure torques."," This means that locally the circulation in the roll is tending to conserve angular momentum, only weakly opposed by any longitudinal fluid pressure torques."
+" In the radiative tachocline with high effective gravity, the strong negative buoyancy forces strong longitudinal fluid pressure torques that are balanced by the Coriolis force on the nearly horizontal motion."," In the radiative tachocline with high effective gravity, the strong negative buoyancy forces strong longitudinal fluid pressure torques that are balanced by the Coriolis force on the nearly horizontal motion."
+ The limited latitudinal extent of the velocity and fluid pressure perturbations in the overshoot case is preferred because that minimizes the stabilizing effect of the perturbations conserving angular momentum., The limited latitudinal extent of the velocity and fluid pressure perturbations in the overshoot case is preferred because that minimizes the stabilizing effect of the perturbations conserving angular momentum.
+C?) ΓΗ ΕΙ). does not vanish there.,"G^2_0 0| _a |0, does not vanish there."
+ Instead. it measures the sea of virtual gluons. the egluon condensate. which defines the difference between the colored svstem and the physical vacuum.," Instead, it measures the sea of virtual gluons, the gluon condensate, which defines the difference between the colored system and the physical vacuum."
+" This anomalous behavior was accounted above through the bag pressure. aud the numerical value of C2 or B can only be determined (""gauged"") empirically. the theory itself being "," This anomalous behavior was accounted above through the bag pressure, and the numerical value of $G^2_0$ or $B$ can only be determined (“gauged”) empirically, the theory itself being scale-invariant."
+In the MIT bag model |12].. one obtains for the radius of a nucleon.," In the MIT bag model \cite{MIT}, one obtains R_N = ( = ( for the radius of a nucleon."
+ For Ry=1 Im. this leads to 42=Gec| /[Im.," For $R_N = 1$ fm, this leads to $4 B = G_0^2 \simeq 1$ $^3$."
+ There are various more refined estimates for the gluon condensate al 7=0. determined by different non-perturbative hadronic inputs [27].. giving a rather wide range of values. G5c1—2 GeV/Im? |[28.29]...," There are various more refined estimates for the gluon condensate at $T=0$, determined by different non-perturbative hadronic inputs \cite{Shif}, giving a rather wide range of values, $G_0^2 \simeq 1 - 2$ $^3$ \cite{SVZ,Leut}."
+" In any case. to recover the correct. να physics, the trace of the energy momentum tensor nist be renormalized |[29].. giving Gi e-ap--. where GF is the expectation value of the glion condensate at temperature T."," In any case, to recover the correct vacuum physics, the trace of the energy momentum tensor must be renormalized \cite{Leut}, giving - 3P = G_0^2 - G_T^2, where $G_T^2$ is the expectation value of the gluon condensate at temperature $T$."
+ In the confinement region T<T.. the latter is found to remain close to GZ [30.31].. leading to e—3D?5-0. until the temperature approaches the deconfinement point.," In the confinement region $T \leq T_c$, the latter is found to remain close to $G_0^2$ \cite{Boyd,Adriano}, leading to $\e - 3P \simeq 0$, until the temperature approaches the deconfinement point."
+" There G5. drops suddenly. i.e.. the gluon condensate melts"" ancl as a consequence. the interaction measure shows a rapid rise. as seen in relinter.."," There $G_T^2$ drops suddenly, i.e., the gluon condensate “melts” and as a consequence, the interaction measure shows a rapid rise, as seen in \\ref{inter}."
+ We (hus find once again (hat the temperature change of the eluon condensate determines (he specific heat of deconfinement: it appears Chat deconfinement corresponds (o ihe melting of the gluon condensate., We thus find once again that the temperature change of the gluon condensate determines the specific heat of deconfinement: it appears that deconfinement corresponds to the melting of the gluon condensate.
+ The second point to note is that the thermocdvnamic observables remain about 10 below their Stefan-Doltzmann values (marked. “SB” in 6)) even at very high temperatures. where the interaction measure becomes very small.," The second point to note is that the thermodynamic observables remain about 10 below their Stefan-Boltzmann values (marked “SB” in \ref{edens}) ) even at very high temperatures, where the interaction measure becomes very small."
+" Such deviations from ideal gas behavior can be modelled in terms of effective thermal masses 17), of quarks and gluons. with macg(D)T. |32.33].."," Such deviations from ideal gas behavior can be modelled in terms of effective `thermal' masses $m_{\rm th}$ of quarks and gluons, with $\m \simeq g(T)~T$ \cite{Golo,Patkos}."
+ Maintaining the next-to-leading order term in mass in the Stefan-Doltzmann form gives for the pressure ο |= e TL =a PTO), Maintaining the next-to-leading order term in mass in the Stefan-Boltzmann form gives for the pressure P = c T^4 ] = c T^4 [1 - a g^2(T)]
+are a significant number of red barred galaxies above the dashed line.,are a significant number of red barred galaxies above the dashed line.
+" Therefore, their result is biased to bluer galaxies."," Therefore, their result is biased to bluer galaxies."
+" Also, they studied the color dependence of foar in the narrow color range of 0.2<g—rS0.7 (see Figure"," Also, they studied the color dependence of $\bfr$ in the narrow color range of $0.2<g-r\lesssim0.7$ (see Figure"
+The nearby galaxy N23 hosts the most luminous steady A-rayv source in the Local Group. N-8.,"The nearby galaxy M33 hosts the most luminous steady X-ray source in the Local Group, X-8."
+ This source. wit Lx~ LO erg i (Lougetal— 1996)). is coincidena to within 5 with the nucleus of the galaxy. (Schuliia&Dregnmau 1995)).," This source, with $L_X\sim$ $^{39}$ erg $^{-1}$ \cite{long}) ), is coincident to within 5” with the nucleus of the galaxy \cite{schulman}) )."
+ Different models were invoked fo AN-8. jueludiug a quiescent ήΔέν (Triuchierietal. lOS: Peresetal. 1989)). a collection of ταν binariey. (IIeruquistetal. 19011) aud a new type of A-ray binary (Gottwald.etal. 1987)).," Different models were invoked for X-8, including a quiescent mini-AGN \cite{trinchieri}; \cite{peres}) ), a collection of X-ray binaries \cite{hernquist}) ) and a new type of X-ray binary \cite{gott}) )."
+ Our ROSAT studies (Dubus 1997)) have shown that X-8 is very steady ou both shor+ and loue tine scales. except for low amplitude (~20%)) variatious which appear modulated on a ~ 106 day period.," Our ROSAT studies \cite{dubus}) ) have shown that X-8 is very steady on both short and long time scales, except for low amplitude $\sim$ ) variations which appear modulated on a $\sim$ 106 day period."
+ This strougly favors a single source explanation for N-8., This strongly favors a single source explanation for X-8.
+" We have interpreted the modulation i N-8 as ""superorbital. simular to that seeu im a unuuber of bright ealactic N-vav binaries wlüch were monitored by e.g. the Vela 5B satellite (Sinale&Lochner 1992))."," We have interpreted the modulation in X-8 as “superorbital”, similar to that seen in a number of bright galactic X-ray binaries which were monitored by e.g. the Vela 5B satellite \cite{smale}) )."
+ N-8 is then likely to be a = LOAL.. black-hole X-ray binary (the high lnass is required to account for the observed hWuuinositv) but with a companion in a nmch shorter orbital period than 106d., X-8 is then likely to be a $\geq$ $M_\odot$ black-hole X-ray binary (the high mass is required to account for the observed luminosity) but with a companion in a much shorter orbital period than 106d.
+" This is supported by the extremely low velocity dispersion of the nucleus which linits the mass of a central black hole in 4133 to <5 &410/ML,. (οποια&AleClure 1993... IKM93)."," This is supported by the extremely low velocity dispersion of the nucleus which limits the mass of a central black hole in M33 to $\leq$ $\times$ $^4$ $_\odot$ \cite{kormendy}, KM93)."
+. This aud the high central stellar density imply that the nucleus is an extremely relaxed. post core-collapse stellar system (comparable for instance to a galactic globular cluster such as MIS).," This and the high central stellar density imply that the nucleus is an extremely relaxed, post core-collapse stellar system (comparable for instance to a galactic globular cluster such as M15)."
+ A sjenificaut umber of stellar collisions/iuteractious could have taken place. eventually leacing to the creation of exotic mteracting binaries (c.e. IIutetal. 1992)).," A significant number of stellar collisions/interactions could have taken place, eventually leading to the creation of exotic interacting binaries (e.g. \cite{hut}) )."
+ The next step toward unravelling the nivsterv of N-8 would be to identity its optical counterpart., The next step toward unravelling the mystery of X-8 would be to identify its optical counterpart.
+" However. even with optimistic Ly/L,,; ratios for either X-ray binaries or ACN the counterpart would only have V~21 compared to a core brightuess of V—11."," However, even with optimistic $L_X/L_{opt}$ ratios for either X-ray binaries or AGN the counterpart would only have $\sim$ 21 compared to a core brightness of $\sim$ 14."
+ But with the optical spectral type of au FP supergiant. the dominance of the ΑΟ visual core cannot extend to UV wavelengths where the hot/flat spectrum of N-S's associated disc ought to be a siguificant contributor.," But with the optical spectral type of an F supergiant, the dominance of the M33 visual core cannot extend to UV wavelengths where the hot/flat spectrum of X-8's associated disc ought to be a significant contributor."
+ This is true despite evidence for a color gradient iu the nucleus (IKM93. Michell&Rich1995.Laueretal. 1998)) sugeesting that a period of recent star formation has taken place and/or that collisions have modified the central star population.," This is true despite evidence for a color gradient in the nucleus (KM93, \cite{mighell,lauer}) ) suggesting that a period of recent star formation has taken place and/or that collisions have modified the central star population."
+ Were we report au attempt to find the counterpart uxiug the UV nuagiug capabilities of IIST., Here we report an attempt to find the counterpart using the UV imaging capabilities of HST.
+ Observations were carricd out with the TST WEPC-2 ou June 12. 1997 uxiug three different filters.," Observations were carried out with the HST WFPC-2 on June 12, 1997 using three different filters."
+ The nucleus was positioned at the ceuter of theCamera (a L:33:51.1. 6 30:39:39. J2000).," The nucleus was positioned at the center of the $\alpha$ 1:33:51.1, $\delta$ 30:39:39, J2000)."
+ During the first orbit two 1200s exposures were made with the FIGOBW CUW filter. Àl191).," During the first orbit two 1200s exposures were made with the F160BW (`UV' filter, $\bar{\lambda} 1491$ )."
+ In the following orbit. two S008 exposures were made with the F300W filter CU filter. A2912À )) aud oue 500s exposure with the F139W filter CD' filter. ÀB00A)).," In the following orbit, two 800s exposures were made with the F300W filter (`U' filter, $\bar{\lambda} 2942$ ) and one 500s exposure with the F439W filter (`B' filter, $\bar{\lambda} 4300$ )."
+ All these exposures were made with the eain setting at 7., All these exposures were made with the gain setting at 7.
+ Tn addition. we have extracted vecalibrated archival data from the Space Telescope European Coordinating Facility (ST-ECE) Archive.," In addition, we have extracted recalibrated archival data from the Space Telescope European Coordinating Facility (ST-ECF) Archive."
+ These data included two LOs exposures with the F555W at eain Li OW filter. AS397A)). wo LOs exposures with the ESLIW at eain 11 CI filter. ÀT92 LÀ)) aud six 300s exposures with the E1012M at eain T CALOT90À)) filters. all centered on the uucleus aud dating roni) September 26-27. 1991.," These data included two 40s exposures with the F555W at gain 14 (`V' filter, $\bar{\lambda} 5397$ ), two 40s exposures with the F814W at gain 14 (`I' filter, $\bar{\lambda} 7924$ ) and six 300s exposures with the F1042M at gain 7 $\bar{\lambda} 10190$ ) filters, all centered on the nucleus and dating from September 26-27, 1994."
+ The archival V aud I data were previously discussed by Laueretal.1998. (L98)., The archival V and I data were previously discussed by \cite{lauer} (L98).
+ Fie., Fig.
+ 1l shows the central region of the reduced UV. U aud D iuages.," 1 shows the central region of the reduced UV, U and B images."
+ The data were reduced using the IIST calibration Xpeline., The data were reduced using the HST calibration pipeline.
+ The signal-to-noise ratio of the only existing FL6OBW flat was quite low., The signal-to-noise ratio of the only existing F160BW flat was quite low.
+ Following advice from the WEPC-2 eroup at STScL we decided to use the F255W flat," Following advice from the WFPC-2 group at STScI, we decided to use the F255W flat"
+The stellar initial mass function (IME) quantifies +ιο distribution of stellar masses in a newly born y.cllay population.,The stellar initial mass function (IMF) quantifies the distribution of stellar masses in a newly born stellar population.
+ Together with the dependency Q stellar evolution on stellar mass and metallicity. as well as the rate at which stars are foriuned iu the Universe. the shape of the IAIF determines the chemical evolution of the Universe aud low its stellar coutent changes with time.," Together with the dependency of stellar evolution on stellar mass and metallicity, as well as the rate at which stars are formed in the Universe, the shape of the IMF determines the chemical evolution of the Universe and how its stellar content changes with time."
+ The shape of the IMF also has important implications for +16 CVOution of star clusters., The shape of the IMF also has important implications for the evolution of star clusters.
+ Thus. kunowiug the VArape of the TIAIF is crucial for a broad variety of astrophysical problems.," Thus, knowing the shape of the IMF is crucial for a broad variety of astrophysical problems."
+ Resolved stellar populations iu the Milkv. Wax and its satellites support the notion that the IAIF does not depend ou the conditions mucder which star formation takes place. but that the stellar masses are distributed according to a single IME known as the INE (227??)..," Resolved stellar populations in the Milky Way and its satellites support the notion that the IMF does not depend on the conditions under which star formation takes place, but that the stellar masses are distributed according to a single IMF known as the IMF \citep{kroupa2001a,kroupa2002a,kumar2008a, bastian2010a}."
+ Ultra conrpact dwarf ealaxies (CCDs) ou the other lac provide evidence for the opposite notion. namely that the IMP. varies aud is top-heavy.," Ultra compact dwarf galaxies (UCDs) on the other hand provide evidence for the opposite notion, namely that the IMF varies and is top-heavy."
+ These CCDs are stellar svsteius that have firs hoe1 discovered in the Fornax galaxy cluster (?).., These UCDs are stellar systems that have first been discovered in the Fornax galaxy cluster \citep{hilker1999a}.
+ They have W-baud huuinosities between 105 aud some 10*L.. but bhalflieht radii of only abou 50 pc or less (C?2)..," They have $V$ -band luminosities between $10^6$ and some $10^7 \ {\rm L}_{\odot}$, but half-light radii of only about 50 pc or less \citep{drinkwater2003a,mieske2008a}."
+" The confined. UCDs are a cistauces where they caunot be resolved iuto stars witi cumeut telescopes. but constrains on their stelar populations cau be set by quantities derived from their iutegrated spectra,"," The confirmed UCDs are at distances where they cannot be resolved into stars with current telescopes, but constrains on their stellar populations can be set by quantities derived from their integrated spectra."
+ Oue such quautity, One such quantity
+LIST observations of the Vela pulsar [iekl were collected between 1997 June aud 2000 Mareh with the WFPC2., observations of the Vela pulsar field were collected between 1997 June and 2000 March with the WFPC2.
+ Several images have been taken through the filters G75I ΤΑ AX= 1536A)). SIAN (A=1995À: AA=1292AÀ)) and 5551 (A=5500A:: AX= 1200A)) with exposure times ranging between 2000 s and 2600 s. These data were originally taken as part of (wo independent programs aimed at the study of the multicolor [αν cistvibution of the pulsar (Mignani Caraveo 2001) and at (he measure of its parallactic displacement (Caraveo et 22001a).," Several images have been taken through the filters $675W$ $\lambda=6717$ $\Delta
+\lambda=1536$ ), $814W$ $\lambda=7995$; $\Delta \lambda =
+1292$ ) and $555W$ $\lambda= 5500$; $\Delta \lambda = 1200$ ) with exposure times ranging between 2000 s and 2600 s. These data were originally taken as part of two independent programs aimed at the study of the multicolor flux distribution of the pulsar (Mignani Caraveo 2001) and at the measure of its parallactic displacement (Caraveo et 2001a)."
+" To achieve (he maximum spatial resolution. in all (he exposures the pulsar was located near ihe center of the Planetary Camera chip. with a pixel size of 07045. and a field. of view of 35""x35""."," To achieve the maximum spatial resolution, in all the exposures the pulsar was located near the center of the Planetary Camera chip, with a pixel size of $0\farcs045$ and a field of view of $35'' \times 35''$."
+ The standardHST pipeline was applied for the image reduction and for the photometric calibrations., The standard pipeline was applied for the image reduction and for the photometric calibrations.
+ Co-aligned exposures were finally combined to filter out cosmic ray hits (see Mignani Caraveo 2001 and Caraveo et 22001a [ον a detailed description of the observations and data reduction)., Co-aligned exposures were finally combined to filter out cosmic ray hits (see Mignani Caraveo 2001 and Caraveo et 2001a for a detailed description of the observations and data reduction).
+ For each observation. the IRAF/STSDAS task was used to obtain a mosaic ol the fou WEPC? images and correct [or the instrumental geometric distortion of the cameras (Casertano Wiges 2001).," For each observation, the IRAF/STSDAS task was used to obtain a mosaic of the four WFPC2 images and correct for the instrumental geometric distortion of the cameras (Casertano Wiggs 2001)."
+ To increase the signal-to-noise ratio. we combined all the available 555MV.EE filter images.," To increase the signal-to-noise ratio, we combined all the available $555W$ filter images."
+ Since the observations al different epochs were performed with different telescope pointines and roll angles. (to superpose (the Irames we had to follow the procedure successfully applied in previous astrometric works. using a common reference frame defined by a grid of reference stars’ coordinates (see Caraveo et 2001a ancl references (herein for farther details).," Since the observations at different epochs were performed with different telescope pointings and roll angles, to superpose the frames we had to follow the procedure successfully applied in previous astrometric works, using a common reference frame defined by a grid of reference stars' coordinates (see Caraveo et 2001a and references therein for further details)."
+ Images were registered within zz0.1 WFC pixels (zz 10 mas). an accuracy adecuate to achieve the &oals of the present investigation. co-aligned and stacked.," Images were registered within $\approx0.1$ WFC pixels $\approx$ 10 mas), an accuracy adequate to achieve the goals of the present investigation, co-aligned and stacked."
+ The Vela pulsar field was observed with theNPT at the La Silla observatory in January 1995 (Nasuti et 1997)., The Vela pulsar field was observed with the at the La Silla observatory in January 1995 (Nasuti et 1997).
+ To achieve the highest sensitivity both at longer and shorter wavelengths. the EMMIEI camera was used in both its Red (EMMI-R) and Blue CEMMI-D) arm modes.," To achieve the highest sensitivity both at longer and shorter wavelengths, the EMMI camera was used in both its Red (EMMI-R) and Blue (EMMI-B) arm modes."
+ The camera is a TEIN 2048x pixels CCD with a pixel size of 07227 (field of view 9/2x8/6 ) and 07337 (field of view 6/2x 622) in the EMMI-R and EMMI-D configuration modes. respectively.," The camera is a TEK $2048\times 2048$ pixels CCD with a pixel size of 27 (field of view $9\farcm2 \times 8\farcm6$ ) and 37 (field of view $6\farcm2 \times 6\farcm2$ ) in the EMMI-R and EMMI-B configuration modes, respectively."
+ Images were collected in the Johnson's U (A=3542À:z AA= 542A). DB (A=4223AÀ5 AA= 941À)). V (A=5426À: AA= 1044A)). and. HR (A= 6410AÀ::AX= 1540À)) bands with exposure times between 15 and 40 minutes.," Images were collected in the Johnson's $U$ $\lambda=3542$; $\Delta \lambda=542$ ), $B$ $\lambda=4223$; $\Delta \lambda=941$ ), $V$ $\lambda=5426$; $\Delta \lambda=1044$ ), and $R$ $\lambda=6410$ ;$\Delta
+\lambda=1540$ ) bands with exposure times between 15 and 40 minutes."
+ Data reduction and photometric calibration were carried out as described in Nasuti et ((1997)., Data reduction and photometric calibration were carried out as described in Nasuti et (1997).
+"where the last equality follows from hydrostatic equilibrium for an isothermal gas (T= 1), that is kpT/mp=d (GM/R)(H/R), with g= GM/R?.","where the last equality follows from hydrostatic equilibrium for an isothermal gas $\Gamma=1$ ), that is $k_{\rm B} T/m_{\rm p} = c_{\rm
+ s}^2 = (GM/R)(H/R)$ , with $g\equiv GM/R^2$ ."
+" We approximate P4 to be the high Reynolds number hydrodynamic drag associated toa cylinder of length L and diameter /, namely (see, e.g. Moreno-Insertis 1986; Landau and Lifshitz 1987) with drag coefficient Ca, of order unity."," We approximate $F_{\rm
+ d}$ to be the high Reynolds number hydrodynamic drag associated toa cylinder of length $L$ and diameter $l$, namely (see, e.g. Moreno-Insertis 1986; Landau and Lifshitz 1987) with drag coefficient $C_{\rm dr}$ of order unity."
+" The work done per unit volume by the net force over a distance H equals the kinetic energy density of the rising structure, H(Fa,—Fy)~ p;U2/2."," The work done per unit volume by the net force over a distance $H$ equals the kinetic energy density of the rising structure, $H(F_{\rm dr}-F_{\rm b})\simeq
+\rho_i U_{\rm b}^2/2$ ."
+" Combining this with Equations (12))-(14)) we obtain Therefore, for all buoyant structures U,(1)>U(l-) since l>l;."," Combining this with Equations \ref{eq:rho_difference}) \ref{eq:Fd}) ) we obtain Therefore, for all buoyant structures $U_{\rm b}(l) \ge U(l_{\rm c})$ since $l\ge l_{\rm c}$."
+" Given that /<Η for all structures that fit in the disk, Up(l)«ος for all | would imply p;/po>2/3—4Car/3z."," Given that $l\le H$ for all structures that fit in the disk, $U_{\rm b}(l) < c_{\rm s}$ for all $l$ would imply $\rho_i/\rho_o > 2/3 - 4C_{\rm dr}/3\pi$."
+ We thus restrict ourselves to this density ratio regime., We thus restrict ourselves to this density ratio regime.
+" Note that our estimate for {7 is itself an upper limit since we consider only a hydrodynamic/ος drag force restricting the buoyant rise, ignoring for example, magnetic tension, c.f. Schramkowski&Torkelsson(2002)."," Note that our estimate for $U_{\rm
+ b}/c_{\rm s}$ is itself an upper limit since we consider only a hydrodynamic drag force restricting the buoyant rise, ignoring for example, magnetic tension, c.f. \citet{ST02}."
+. Whether the dynamics allows densities below the above upper limit remains an open question., Whether the dynamics allows densities below the above upper limit remains an open question.
+" However, even if the structures were initially able to move faster than c,, we would expect shocks and the associated dissipation to slow the motion via additional drag."," However, even if the structures were initially able to move faster than $c_s$, we would expect shocks and the associated dissipation to slow the motion via additional drag."
+ Vishniac (1995) estimated a buoyant rise time acs.," Vishniac (1995) estimated a buoyant rise time $U_{\rm b}\sim \alpha
+c_{\rm s}$ ."
+" This value is low compared to the upper limit c, we considered.", This value is low compared to the upper limit $c_{\rm s}$ we considered.
+" Had we used Uy~«oc, then the value of f, in Equation (17)) would dramatically increase; highlighting the importance of large scale field for coronal dynamics even more."," Had we used $U_{\rm b} \sim \alpha c_{\rm s}$, then the value of $f_{\rm s}$ in Equation \ref{eq:fs_low_lim}) ) would dramatically increase; highlighting the importance of large scale field for coronal dynamics even more."
+ Some characteristic values of f; resulting from this smaller buoyancy speed are shown as dashed lines in Figure 2.., Some characteristic values of $f_{\rm s}$ resulting from this smaller buoyancy speed are shown as dashed lines in Figure \ref{fig:fs_vs_q}.
+" Setting |=|, in Equation (15)), and using Equation (2)), gives a 5' order equation for |./H as a function of Car, po, and a."," Setting $l=l_{\rm c}$ in Equation \ref{eq:Ub_to_cs}) ), and using Equation \ref{eq:lc}) ), gives a $5^{\rm th}$ order equation for $l_{\rm c}/H$ as a function of $C_{\rm dr}$, $\rho_i/\rho_o$, and $\alpha$."
+" The physical solution to this equation is shown in pi/Figure 1 forarange in drag coefficients, 10-3<Ca2, and for three values of a={0.5,0.1,0.01} as a function of the density ratio ϱι/βρο."," The physical solution to this equation is shown in Figure \ref{fig:lc_to_H} for a range in drag coefficients, $10^{-3}\le
+C_{\rm dr} \le 2$, and for three values of $\alpha=\{0.5, 0.1, 0.01\}$ as a function of the density ratio $\rho_{\rm i}/\rho_o$ ."
+ The characteristic turbulent scale EEa1/?H is also shown as horizontal lines for each value of a., The characteristic turbulent scale $l_{\rm t} \equiv \alpha^{1/2}H$ is also shown as horizontal lines for each value of $\alpha$.
+" Since ἰς is the minimum scale for buoyant rise, the fact that {ο>1, highlights the importance of large scale fields."," Since $l_{\rm c}$ is the minimum scale for buoyant rise, the fact that $l_{\rm c}> l_{\rm t}$ highlights the importance of large scale fields."
+" For the regime on the plot where /,>H, a buoyant corona cannot arise from fields produced internally to the disk."," For the regime on the plot where $l_{\rm c}>H$, a buoyant corona cannot arise from fields produced internally to the disk."
+" We can set two constraints on the minimum fraction f, of magnetic energy that the disk must produce in fields with scales 1>I, for a given coronal to bolometric emission fraction g.", We can set two constraints on the minimum fraction $f_{\rm s}$ of magnetic energy that the disk must produce in fields with scales $l>l_{\rm c}$ for a given coronal to bolometric emission fraction $q$.
+" The more stringent bound relies on the fact that U,(l)« for |«H;a less severe limit is obtained requiringUy(H)U,(H)<cs, which is satisfied as long as the density ratio p;/po does not fall below the lower limit discussed above."," The more stringent bound relies on the fact that $U_{\rm b}(l) < U_{\rm
+ b}(H)$ for $l<H$; a less severe limit is obtained requiring$U_{\rm
+ b}(H)<c_{\rm s}$, which is satisfied as long as the density ratio $\rho_i/\rho_o$ does not fall below the lower limit discussed above."
+" Applying these conditions to Equation (10)) we obtain Then,for an observationally inferred value of coronal to bolometric flux, we can obtain a lower limit on the fraction of"," Applying these conditions to Equation \ref{eq:Dc_to_Dd}) ) we obtain Then,for an observationally inferred value of coronal to bolometric flux, we can obtain a lower limit on the fraction of"
+The vertical structure of an exoplanetary atmosphere can. be derived from transmission and emission spectroscopy when a planet transits its host star.,The vertical structure of an exoplanetary atmosphere can be derived from transmission and emission spectroscopy when a planet transits its host star.
+ The atmospheric properties of HD 189733b have been extensively investigated using transit spectroscopy observations. from. both space and. ground-based telescopes., The atmospheric properties of HD 189733b have been extensively investigated using transit spectroscopy observations from both space and ground-based telescopes.
+ Transmission spectra acquired during primary transits. where light from the host star is filtered through the upper layers of an exoplanetary atmosphere. have been measured in the near-. mid- and far-infrared (NIR. MIR and FIR. respectively) by both the LLubble Space Telescope LEST) (Swainetal.2008:Singal.2009). and 9piízer telescope (Tinettietal.2007a:Beaulieuetal.2008:Désert 2009).," Transmission spectra acquired during primary transits, where light from the host star is filtered through the upper layers of an exoplanetary atmosphere, have been measured in the near-, mid- and far-infrared (NIR, MIR and FIR, respectively) by both the $Hubble$ $Space$ $Telescope$ $HST$ ) \citep{swa08,sin09} and $Spitzer$ telescope \citep{tin07a,bea08,des09}."
+.. Spi/zer mid- and far-infrared (MIR and FIR) broadband photometry was used to deduce a high abundance of ΠΟΟ in the terminator region of this planet2009)., $Spitzer$ mid- and far-infrared (MIR and FIR) broadband photometry was used to deduce a high abundance of $_{2}$ O in the terminator region of this planet.
+". In addition. a strong absorption due to CH, was reported in the near-infrared (NIR) using the NICMOS (Near Infrared Camera and Multi-Object Spectrometer) spectrophotometry on the //ST (Swainetal.2008)."," In addition, a strong absorption due to $_{4}$ was reported in the near-infrared (NIR) using the NICMOS (Near Infrared Camera and Multi-Object Spectrometer) spectrophotometry on the $HST$ \citep{swa08}."
+". However. the accuracy of transit spectroscopy and photometry has been called into question (seethereviewbySeager&Deming 2010)3.. and most recently Gibsonetal.(2011) claimed that instrumental systematics limit our ability to constrain atmospheric CH, from these //S7 7/NICMOS data as well as CO: and CO."," However, the accuracy of transit spectroscopy and photometry has been called into question \citep[see the review by][]{sea10}, and most recently \citet{gib11} claimed that instrumental systematics limit our ability to constrain atmospheric $_{4}$ from these $HST$ /NICMOS data as well as $_{2}$ and CO."
+ A spectral feature at 4.5 j/im can be explained by either of the carbon bearing molecules (Ου or CO. although it is still unclear which due to the degeneracy of the solution (Désertetal.2009:Fortneyetal. 2010).," A spectral feature at 4.5 $\mu$ m can be explained by either of the carbon bearing molecules $_{2}$ or CO, although it is still unclear which due to the degeneracy of the solution \citep{des09,for10}."
+ At visible wavelengths. strong lines of the alkali metal sodium were reported using observations with the high resolution ground spectrograph of the Llobby—Lberly Telescope (Redfieldetal.2008).. while featureless spectra detected by the LEST/ACS (Pontetal.2008) and // ST/STIS (Singetal.2011) are thought to be caused by thick atmospheric hazes.," At visible wavelengths, strong lines of the alkali metal sodium were reported using observations with the high resolution ground spectrograph of the $\it{Hobby}$ $\it{Eberly}$ Telescope \citep{red08}, while featureless spectra detected by the $HST$ /ACS \citep{pon08} and $HST$ /STIS \citep{sin11} are thought to be caused by thick atmospheric hazes."
+ Observations of secondary transit spectra of HD 1507330 by Spilzer broadband photometry. when emission spectra from," Observations of secondary transit spectra of HD 189733b by $Spitzer$ broadband photometry, when emission spectra from"
+"below, the high velocity dispersions of the EBS and Hydra I may be due to significant heating by encounters with the dark matter subhalos 2009) or massive star clusters or molecular clouds(Carlberg in the disk.","below, the high velocity dispersions of the EBS and Hydra I may be due to significant heating by encounters with the dark matter subhalos \citep{carlberg2009} or massive star clusters or molecular clouds in the disk."
+" Mindful of the fact that tidal streams do not precisely trace the orbits of their progenitors (Odenkirchenetal.2009;Eyre&Binney2009, 2011), we nevertheless estimate the orbit of thestream stars to determine whether the EBS might be related to any of the other streams in Figure 1."," Mindful of the fact that tidal streams do not precisely trace the orbits of their progenitors \citep{odenkirchen2009, eyre2009, eyre2011}, we nevertheless estimate the orbit of the stars to determine whether the EBS might be related to any of the other streams in Figure 1."
+" We do this using the Galactic model of Allen&Santillan to compute trial orbit integrations, and matching (1991)these orbit integrations with the measured positions, distances, and velocities of the stream in a least-squares sense."," We do this using the Galactic model of \citet{allen91} to compute trial orbit integrations, and matching these orbit integrations with the measured positions, distances, and velocities of the stream in a least-squares sense."
+" We integrate orbits over a grid of possible radial velocities and proper motions, using the IDL AMOEBA downhill simplex procedure to find the minimum x? at each grid point."," We integrate orbits over a grid of possible radial velocities and proper motions, using the IDL AMOEBA downhill simplex procedure to find the minimum $\chi^2$ at each grid point."
+" The grid points are separated by 1 km s~! in radial velocity, and 0.02 mas yr-! in each component of proper motion."," The grid points are separated by 1 km $^{-1}$ in radial velocity, and 0.02 mas $^{-1}$ in each component of proper motion."
+ We fit to 17 normal points chosen to lie along the estimated centerline of the stream., We fit to 17 normal points chosen to lie along the estimated centerline of the stream.
+" We use a solar Galactocentric distance of 8.5 kpc, and stream distances and velocities as given above."," We use a solar Galactocentric distance of 8.5 kpc, and stream distances and velocities as given above."
+" We adopt velocity uncertainties for the measurements at B-7/PCI-8/PCII-20 and B-8/PCI-9/PCII-21 of 3 km s! (Schlaufman, private positional uncertainties of 0.2?, and distance communication),uncertainties as given above."," We adopt velocity uncertainties for the measurements at B-7/PCI-8/PCII-20 and B-8/PCI-9/PCII-21 of 3 km $^{-1}$ (Schlaufman, private communication), positional uncertainties of $0.2\arcdeg$, and distance uncertainties as given above."
+" We have attempted to measure the proper motions along the stream using the proper motions provided in DR7 (Munnetal.2004,"," We have attempted to measure the proper motions along the stream using the proper motions provided in DR7 \citep{munn2004, munn2008}."
+" Unfortunately, the relatively large uncertainties and 2008)..severe contamination by field stars conspire to wash out any obvious signal due to the stream."," Unfortunately, the relatively large uncertainties and severe contamination by field stars conspire to wash out any obvious signal due to the stream."
+ We consequently leave the proper motions as free parameters in the fit., We consequently leave the proper motions as free parameters in the fit.
+ Tighter constrains on the orbit will have to await the acqusition of more accurate proper motions with Gaia or the Large Synoptic Survey Telescope., Tighter constrains on the orbit will have to await the acqusition of more accurate proper motions with Gaia or the Large Synoptic Survey Telescope.
+ Figure 7 shows the normal points used to fit the postions along the EBS stream., Figure 7 shows the normal points used to fit the postions along the EBS stream.
+ Also shown are projections of the best-fit orbits on the SDSS DR7 footprint., Also shown are projections of the best-fit orbits on the SDSS DR7 footprint.
+" Both the prograde and retrograde orbit models predict that the stream passes between 1 and 2 kpc of the current position of the sun, near the northGalactic pole in projection."," Both the prograde and retrograde orbit models predict that the stream passes between 1 and 2 kpc of the current position of the sun, near the northGalactic pole in projection."
+" At this distance, the stream would be more than 5° across, and moving between three and four times faster than at apogalacticon."," At this distance, the stream would be more than $5\arcdeg$ across, and moving between three and four times faster than at apogalacticon."
+" If the stellar stream extends along these portions of the orbits, and if the number of stars stripped from the progenitor per unit time was roughly constant over the lifetime of the stream, then we would expect the surface density near the north Galactic pole to be ~15—20 times less than it is in Figure 1."," If the stellar stream extends along these portions of the orbits, and if the number of stars stripped from the progenitor per unit time was roughly constant over the lifetime of the stream, then we would expect the surface density near the north Galactic pole to be $\sim 15 - 20$ times less than it is in Figure 1."
+" Since we expect the rate of tidal stripping to accelerate as the mass of the progenitor is diminished over time, the relative number of stream stars we might expect to see far from the progenitor would be reduced further still."," Since we expect the rate of tidal stripping to accelerate as the mass of the progenitor is diminished over time, the relative number of stream stars we might expect to see far from the progenitor would be reduced further still."
+" If the orbit is retrograde, then the star count signal of the EBS would likely be buried within the much more populous Sagittarius stream."," If the orbit is retrograde, then the star count signal of the EBS would likely be buried within the much more populous Sagittarius stream."
+" If there are EBS stars passing near the sun, they will more easily be found in velocity and proper motion surveys (e.g. RAVE, Gaia, LSST)."," If there are EBS stars passing near the sun, they will more easily be found in velocity and proper motion surveys (e.g. RAVE, Gaia, LSST)."
+ Figure 8 shows orbit integrations that correspond to the best-fit parameters for both prograde and retrograde models., Figure 8 shows orbit integrations that correspond to the best-fit parameters for both prograde and retrograde models.
+" While the formal x? for the retrograde model is less than that for the prograde model, the disagreement betweenthe predicted velocity gradient along the stream (3.1 km s! deg!) and the measurements at B-7/PCI-8/PCII-20 and B-8/PCI-9/PCII-21 (-4.8 km s! 4ος1) is somewhat larger (and of the opposite sign) than for the prograde model (-10.7 km s! deg-!)."," While the formal $\chi^2$ for the retrograde model is less than that for the prograde model, the disagreement betweenthe predicted velocity gradient along the stream (3.1 km $^{-1}$ $^{-1}$ ) and the measurements at B-7/PCI-8/PCII-20 and B-8/PCI-9/PCII-21 (-4.8 km $^{-1}$ $^{-1}$ ) is somewhat larger (and of the opposite sign) than for the prograde model (-10.7 km $^{-1}$ $^{-1}$ )."
+ Additional velocity measurements at different positions along the stream and/or proper motion measurements will be required to resolve the ambiguity., Additional velocity measurements at different positions along the stream and/or proper motion measurements will be required to resolve the ambiguity.
+" The proper motions predicted at the position of Hydra I are µαcosó=—0.15 mas yr!, us=—2.67 mas yr-! for the prograde orbit, and j44cosó=+0.65 mas yr-!, jig=—5.08 mas yr! for the retrograde orbit."," The proper motions predicted at the position of Hydra I are $\mu_\alpha \cos{\delta} = -0.15$ mas $^{-1}$ , $\mu_\delta = -2.67$ mas $^{-1}$ for the prograde orbit, and $\mu_\alpha \cos{\delta} =
+ +0.65$ mas $^{-1}$, $\mu_\delta = -5.08$ mas $^{-1}$ for the retrograde orbit."
+" These values are of the same magnitude or less than the typical DR7 uncertainties (e4 mas yr—1), so it is perhaps not surprizing that we have been unable to identify a clear stream signature in the proper motion data."," These values are of the same magnitude or less than the typical DR7 uncertainties $\approx 4$ mas ${-1}$ ), so it is perhaps not surprizing that we have been unable to identify a clear stream signature in the proper motion data."
+" The prograde orbit model predicts apogalacticon R4=16.5+0.1 kpc, perigalacticon Rp=3.073% kpc, eccentricity e=0.697002, and inclination i=17° 0.4°, where the uncertainties correspond to the joint confidence interval."," The prograde orbit model predicts apogalacticon $_A = 16.5 \pm 0.1$ kpc, perigalacticon $_P =3.0^{+0.7}_{-0.3}$ kpc, eccentricity $\epsilon = 0.69^{+0.02}_{-0.05}$, and inclination $i = 17\arcdeg \pm
+0.4\arcdeg$ , where the uncertainties correspond to the joint confidence interval."
+" For the retrograde orbit, the parameters are R4=17.740.4 kpc, Rp=18103 kpc, c=0.8703, and i= 19.9°+0.7°."," For the retrograde orbit, the parameters are $_A
+= 17.7 \pm 0.4$ kpc, $_P = 1.8^{+0.4}_{-0.2}$ kpc, $\epsilon =
+0.8^{+0.02}_{-0.03}$, and $i = 19.9\arcdeg \pm 0.7\arcdeg$ ."
+" In either case, the EBS appears to be on a fairly eccentric orbit."," In either case, the EBS appears to be on a fairly eccentric orbit."
+ This eccentricity is forced entirely by the curvature of the stream at its northern end., This eccentricity is forced entirely by the curvature of the stream at its northern end.
+" If we give zero weight to the B-7/PCI-8/PCIL-20 and B-8/PCI-9/PCII-21 velocities, the best-fitting orbit still predicts an eccentricity of ezz 0.8."," If we give zero weight to the B-7/PCI-8/PCII-20 and B-8/PCI-9/PCII-21 velocities, the best-fitting orbit still predicts an eccentricity of $\epsilon \approx 0.8$ ."
+" The stream is evidently not associated with either the ACS or the Monoceros Ring, both of whichhave been determined to be on very nearlycircular orbits (Penarrubiaet 2008).."," The stream is evidently not associated with either the ACS or the Monoceros Ring, both of whichhave been determined to be on very nearlycircular orbits \citep{penarrubia2005, grill2008}. ."
+" The combination of inclination and eccentricity takes the stream into the inner, non-spherical part of the Galactic"," The combination of inclination and eccentricity takes the stream into the inner, non-spherical part of the Galactic"
+Bars are common in disk galaxies. and are thought to jiwe an important impact on the evolution of galaxies hrough their ability to transfer angular momentum in both he barvonic and dark matter components of the galaxy (Combes&Sanders1981:Weinberg1985:Debattistawood2000:Berentzenetal. 2006).,"Bars are common in disk galaxies, and are thought to have an important impact on the evolution of galaxies through their ability to transfer angular momentum in both the baryonic and dark matter components of the galaxy \citep{CS81,W85,DS00,B06}."
+. Bars are cllicicnt Sell-at driving gas inwards. perhaps sparking central star formation thus help to grow a central 2005:xilee (c.g.Shethforaal.review2005)..Ixo," Bars are efficient at driving gas inwards, perhaps sparking central star formation, and thus help to grow a central bulge \citep[e.g. for a review][]{KK04}."
+rmendyand& Bars may also feed a central black hole (as per IxennicuttShlosman2004)..ctal.1089.1090. and see Jogee2006. [or a recent review). but so zw no correlation has been found between AGN activity and xw [fraction in galaxies forming stars (c.g. Lloetal.1997 and as recently discussed by Haoetal. 2009).," Bars may also feed a central black hole (as per \citealt{S89,S90} and see \citealt{J06} for a recent review), but so far no correlation has been found between AGN activity and bar fraction in galaxies forming stars (e.g. \citealt{H97} and as recently discussed by \citealt {H09}) )."
+ Early visual inspectionof spiral galaxies in catalogues ike the RCS (deVaucouleursetal...1991). an eaveoptical xw fraction of fi;~0.25 0.3. rising to if weaker," Early visual inspectionof spiral galaxies in catalogues like the RC3 \citep{RC3} gave an optical bar fraction of $f_{\rm bar} \sim 0.25-0.3$ , rising to if weaker"
+spectrin (see Fig. 119) ,spectrum (see Fig. \ref{fig:xfit}) )
+and found that the ft improved sieuificautlv., and found that the fit improved significantly.
+ Thus. we cousider our previous finding of at least two spectral coniponeuts in Holl N-1 coufirmed.," Thus, we consider our previous finding of at least two spectral components in HoII X-1 confirmed."
+ The results of the power-law plus thermal fits [fit (B)| are also shown in Table 6.., The results of the power-law plus thermal fits [fit (B)] are also shown in Table \ref{tab:xfit}.
+ The thermal component is clearly soft (sco Fig. 15)), The thermal component is clearly soft (see Fig. \ref{fig:uxfit}) )
+ aud coutributes about 18 to the total 38.0 keV flux., and contributes about 18 to the total 0.3-8.0 keV flux.
+ A power-law fit to the soft component is ruled out u=566 for 176 d.o.f.)., A power-law fit to the soft component is ruled out $\chi^{2}=566$ for 476 d.o.f.).
+ Au absorbed power-law and a inultieolor disk black ον (1n0del € πι Table 6)) gives a good fit as well., An absorbed power-law and a multicolor disk black body (model C in Table \ref{tab:xfit}) ) gives a good fit as well.
+ This neaus that the soft component can be as well described » à MCD model., This means that the soft component can be as well described by a MCD model.
+ The imuer disk temperature is quite ow GT~0.16 keV). which indicate a coo accretion disk.," The inner disk temperature is quite low $kT\sim0.16$ keV), which indicate a cool accretion disk."
+ Soft components with a cool accretion disks may indicate iutermediate-nass black holes with a black hole mass of Afpyzz 10M; (Miller et al; 2003))., Soft components with a cool accretion disks may indicate intermediate-mass black holes with a black hole mass of $M_{BH}\simeq10^{3}$ $_{\sun}$ (Miller et al. \cite{Mil03}) ).
+" The excess of cold absorption (Nyo~1.6«1024 7) over the galacetie value (Ngδες107"" 2) is confirmed bv all fits (A to D).", The excess of cold absorption $N_{\rm H}\sim 1.6\times 10^{21}$ $^{-2}$ ) over the galactic value $N_{\rm H}\sim 3.4\times 10^{20}$ $^{-2}$ ) is confirmed by all fits (A to D).
+ The intrinsic cold absorption is iu good agrecineut with recent VLA observations of Wolubere IL., The intrinsic cold absorption is in good agreement with recent VLA observations of Holmberg II.
+ Bureau Cariguau (20023) showed that the compact IT IT region #770 is located inside one of the regions with the largest intrinsic hydrogen col density 1073 2)., Bureau Carignan \cite{Bu02}) ) showed that the compact H II region 70 is located inside one of the regions with the largest intrinsic hydrogen column density $\times 10^{21}$ $^{-2}$ ).
+ A power-law and a thermal thin plasima model using a solar metal abundance |fit (D)| also eives a resonable fit., A power-law and a thermal thin plasma model using a solar metal abundance [fit (D)] also gives a resonable fit.
+ The fraction of the 0.3-8.0 keV hermal thin flux (auckal) to the total intrinsic flux is about 9%., The fraction of the 0.3-8.0 keV thermal thin flux (mekal) to the total intrinsic flux is about 9.
+. The hick plasma thermal emission ybody) is assiuuued to be originating from the accretion disk around the black hole., The thick plasma thermal emission (bbody) is assumed to be originating from the accretion disk around the black hole.
+" ""Therefore. the thick plasima thermal ciission is spatially rot resolved."," Therefore, the thick plasma thermal emission is spatially not resolved."
+ The spatial analysis of the ACTS-S nuage sugeests a possible extended component with a fraction of less than 15 to the total 0.3-8.0 keV Hux. which would be iu aereemoent with the flux fraction resulting from the thermal thin plasima.," The spatial analysis of the ACIS-S image suggests a possible extended component with a fraction of less than 15 to the total 0.3-8.0 keV flux, which would be in agreement with the flux fraction resulting from the thermal thin plasma."
+ This would favor fit (D)., This would favor fit (D).
+ A power-law model and a combination of a thick aud a thin thermal plasina model (E) gives ouly a mareinal iuproveineut of the fit., A power-law model and a combination of a thick and a thin thermal plasma model (E) gives only a marginal improvement of the fit.
+ The iutriusic dux iu the 0.3-8.0, The intrinsic flux in the 0.3-8.0
+"(0.17 with a 1-day mean flux binning using data, and 0.26 with a 0.5-day binning using data) found in the most X-ray variable narrow-line Seyfert 1 galaxy IRAS 13244-3809 by Gaskell(2004)., the only other known case of an extragalactic source exhibiting lognormal ray flux variations.","(0.17 with a 1-day mean flux binning using data, and 0.26 with a 0.5-day binning using data) found in the most X-ray variable narrow-line Seyfert 1 galaxy IRAS $-$ 3809 by \cite{gas04}, the only other known case of an extragalactic source exhibiting lognormal X-ray flux variations."
+ The width in Fig., The width in Fig.
+" 2 of σ=0.26 is however smaller (0.44dex and 0.54dex for the two different data sets mentioned above) than in IRAS 13244-3809, meaning that the fluctuations are smaller in this BL Lac than in the Seyfert galaxy."," \ref{fig2}
+ of $\sigma=0.26$ is however smaller $0.44\,{\rm dex}$ and $0.54\,{\rm dex}$ for the two different data sets mentioned above) than in IRAS $-$ 3809, meaning that the fluctuations are smaller in this BL Lac than in the Seyfert galaxy."
+ A linear correlation between the excess rms and the average flux was also found in the Seyfert 1 galaxy Mrk 766 (Vaughanetal. 2003))., A linear correlation between the excess rms and the average flux was also found in the Seyfert 1 galaxy Mrk 766 \cite{vau03}) ).
+ The physical process responsible for X-ray emission in Seyfert 1 galaxies is probably thermal emission from the accretion disk., The physical process responsible for X-ray emission in Seyfert 1 galaxies is probably thermal emission from the accretion disk.
+" We also note that a linear flux-rms relationship was found for Galactic accreting sources such as X-ray binaries Cyg X-1 and SAX J 1808.4—3658 by Uttley&M*Hardy(2001)., of proportionality factors of 0.33 and 0.31, respectively."," We also note that a linear flux-rms relationship was found for Galactic accreting sources such as X-ray binaries Cyg X-1 and SAX J $-$ 3658 by \cite{utt01}, of proportionality factors of 0.33 and 0.31, respectively."
+ A specific observation of Cyg X-1 exhibited a remarkable lognormal fit to its X-ray flux distribution (Uttleyetal. 2005))., A specific observation of Cyg X-1 exhibited a remarkable lognormal fit to its X-ray flux distribution \cite{utt05}) ).
+ The origin of the linear flux-rms was interpreted by Uttley&M*Hardy(2001) as originating from the subdivision of magnetic reconnection energy release as an avalanche occuring on large scales in the corona., The origin of the linear flux-rms was interpreted by \cite{utt01} as originating from the subdivision of magnetic reconnection energy release as an avalanche occuring on large scales in the corona.
+ The same authors note that the radius-dependent fluctuations in the mass accretion rate in the model of Lyubarskiietal.(1997) can also explain this relationship., The same authors note that the radius-dependent fluctuations in the mass accretion rate in the model of \cite{lyu97} can also explain this relationship.
+" It seems possible that the photon breeding mechanism (Stern&Poutanen2008,, and references therein), which generates highly non-linear light curves in relativistic jets and operates in sites that are not necessarily directly connected to accretion properties, could be an alternative explanation of lognormal blazar variability, but this remains to be investigated since the statistical properties of the photon luminosity from this model are not known at this point."," It seems possible that the photon breeding mechanism \cite{ste08}, and references therein), which generates highly non-linear light curves in relativistic jets and operates in sites that are not necessarily directly connected to accretion properties, could be an alternative explanation of lognormal blazar variability, but this remains to be investigated since the statistical properties of the photon luminosity from this model are not known at this point."
+" These striking similarities suggest that the variability mechanism causing the X-ray modulation in BL Lac, outside of flaring events, is independent of the emission mechanism, and originates in the disk, rather than in the jet, provided these lognormal fluctuations are indeed present in the disk."," These striking similarities suggest that the variability mechanism causing the X-ray modulation in BL Lac, outside of flaring events, is independent of the emission mechanism, and originates in the disk, rather than in the jet, provided these lognormal fluctuations are indeed present in the disk."
+" The nonthermal emission from BL Lac is usually modeled by the synchrotron self-Compton scenario (SSC, see e.g., Maraschi,Ghisellini&Celotti1992,Bicknell&Wagner 2002)), where relativistic electrons cool either via the synchrotron or the inverse Compton mechanism."," The nonthermal emission from BL Lac is usually modeled by the synchrotron self-Compton scenario (SSC, see e.g., \cite{mar92,bic02}) ), where relativistic electrons cool either via the synchrotron or the inverse Compton mechanism."
+" The comptonized photons are assumed to be predominantly the synchrotron photons themselves, rather than photons coming from the outside the emission region as in the external Compton model (EC, see e.g. Sikoraetal. 1994)), given the absence of thermal signatures in the emission spectrum of BL Lac."," The comptonized photons are assumed to be predominantly the synchrotron photons themselves, rather than photons coming from the outside the emission region as in the external Compton model (EC, see e.g. \cite{sik94}) ), given the absence of thermal signatures in the emission spectrum of BL Lac."
+" However, these mechanisms do not provide a way to generate lognormal variability if one of the parameters affecting the predicted flux (e.g., magnetic field, Doppler factor, injection rate) does not have a lognormal time-dependence itself."," However, these mechanisms do not provide a way to generate lognormal variability if one of the parameters affecting the predicted flux (e.g., magnetic field, Doppler factor, injection rate) does not have a lognormal time-dependence itself."
+" If fluctuations in the disk are indeed the origin of the lognormal radiation in BL Lac, then the most appropriate hypothesis would be one in which lognormal fluctuations in the accretion rate yield an injection rate with similar properties."," If fluctuations in the disk are indeed the origin of the lognormal radiation in BL Lac, then the most appropriate hypothesis would be one in which lognormal fluctuations in the accretion rate yield an injection rate with similar properties."
+" If however external photons are at least partly responsible for the cooling of the X-ray emitting electrons, then lognormal fluctuations in the disk luminosity would have a similar result."," If however external photons are at least partly responsible for the cooling of the X-ray emitting electrons, then lognormal fluctuations in the disk luminosity would have a similar result."
+" An ambitious observation programme targeting other X-ray bright blazars, over similar time scales, and with a similar sampling quality, would allow us to search for lognormal variability, and perhaps solve an important piece in the puzzle about blazar variability."," An ambitious observation programme targeting other X-ray bright blazars, over similar time scales, and with a similar sampling quality, would allow us to search for lognormal variability, and perhaps solve an important piece in the puzzle about blazar variability."
+" We thank the anonymous referee for the careful reading and the constructive remarks, which have improved this report."," We thank the anonymous referee for the careful reading and the constructive remarks, which have improved this report."
+ BG gratefully acknowledges comments from D. Barret and G. Dubus on an early draft., BG gratefully acknowledges comments from D. Barret and G. Dubus on an early draft.
+After and.. the nost abundaut nucleus i a M«M. white dwarf (WD) interior is?2Ne.,"After and, the most abundant nucleus in a $M<M_\sun$ white dwarf (WD) interior is."
+. The reason for this starts with the physics of the CNO evcle in the hydrogen burning phase of a stars lite., The reason for this starts with the physics of the CNO cycle in the hydrogen burning phase of a star's life.
+ The slowest step in the CNO evele is the proton capture onto and almost all of the CNO catalyst uncle eud up as at the completion of burnius., The slowest step in the CNO cycle is the proton capture onto and almost all of the CNO catalyst nuclei end up as at the completion of burning.
+ The (aΙΕ (313150. (α.1 reaction. sequeuce during the subsequent helimm burning then processes all Πιο intoΝ," The $(\alpha\,,\gamma)$ $(\beta^+)$ $(\alpha,\gamma)$ reaction sequence during the subsequent helium burning then processes all the into."
+"ο, This results in a mass fraction of Noo29ZongO.02 for receutly formed WDs of AJ<AL. (sec. for example. Uinedaetal. (1999))) made from stars =OAL..."," This results in a mass fraction of $X_{22}\approx Z_{\mathrm{CNO}}\approx 0.02$ for recently formed WDs of $M<M_\sun$ (see, for example, \citet{um99}) ) made from stars of initial mass $\lesssim 6 \msun$."
+ The possibility that could play a role iu the enuergetics of a coolne WD was first noted bv Is«cxuoetal(1991)., The possibility that could play a role in the energetics of a cooling WD was first noted by \citet{isern91}.
+ Thev discussed the possibility of gravitational cucrey release i£ pliase separates at core crystallization., They discussed the possibility of gravitational energy release if phase separates at core crystallization.
+ This was followed by more detailed studies (δα&VanHorn1992:Ogatactal.1993:Seeretainet1991:Seeretain1996) which differed in their couclusious about the final state ounce crystallization was complete: ranging from all the iu the stellar center to most of it havine an unchanged profile.," This was followed by more detailed studies \citep{xv92, ogata93, seg94, seg96} which differed in their conclusions about the final state once crystallization was complete; ranging from all the in the stellar center to most of it having an unchanged profile."
+ Bildsten&Tall(2001) (hereafter BIOL) examined another aspect of which had been previously overlooked (except for a brief mcution in Bravo et al., \citet{bild01} (hereafter BH01) examined another aspect of which had been previously overlooked (except for a brief mention in Bravo et al.
+ 1992). namely its ability to sux through the liquid WD interior.," 1992), namely its ability to sink through the liquid WD interior."
+" As discussed in BITOL. there is an upward pointing electric field. of magnitude (Ez2i,g ta the degenerate WD interior. where g is the local gravitational acceleration iid i, is the proton mass."," As discussed in BH01, there is an upward pointing electric field of magnitude $eE\approx 2m_p g$ in the degenerate WD interior, where $g$ is the local gravitational acceleration and $m_p$ is the proton mass."
+" The net force ou is therefore F— 22,gP|Wek νο the diffusion of inward."," The net force on is therefore $\vec{F}=- 22m_p g \rhat + 10eE \rhat=- 2m_pg \rhat$, biasing the diffusion of inward."
+ It is the excess ueutrons of the nucleus. (velative to the predominant 4d=27 nuclei. A being the atomic mass umber and Z the charge) that aalke it special in this regard.," It is the excess neutrons of the nucleus (relative to the predominant $A=2Z$ nuclei, $A$ being the atomic mass number and $Z$ the charge) that make it special in this regard."
+ BIOL found that the eravitational enerev released by sedimenting is comparable to the thermal couteut of the WD at an age of afew Cyrs., BH01 found that the gravitational energy released by sedimenting is comparable to the thermal content of the WD at an age of a few Gyrs.
+ This raised the question of whether could settle fast enough to significantly affect the WD's cooling sienificantly., This raised the question of whether could settle fast enough to significantly affect the WD's cooling significantly.
+ The rate at which sinks im the liquid C/O WD depends ou the diffusion coefficieut. D.," The rate at which sinks in the liquid C/O WD depends on the diffusion coefficient, $D$."
+ The conditions there are ναν non-leal. as the average Coulomb energy for the ious is comparable to the thermal encrgy. or P—(Ze/akT»1. where a?—Anuop.," The conditions there are very non-ideal, as the average Coulomb energy for the ions is comparable to the thermal energy, or $\Gamma\equiv (Ze)^2/ a kT>1$, where $a^3=3Am_p/4\pi \rho$."
+" As discussed by Paquetteetal.(1986)... there is substautial (factors of many) uncertaiutv in the diffusion coefficients in these liquid regimes, as the familiar notions of mean free path lose their meaning."," As discussed by \citet{paq86}, there is substantial (factors of many) uncertainty in the diffusion coefficients in these liquid regimes, as the familiar notions of mean free path lose their meaning."
+ Iun the absence of a definitive calculation of D for this situation. DIIOl proceeded by estimating D for in a C/O plasma by the κοdiffusion coefficient. Dy. of the classical oue-componcut plasma (OCP).," In the absence of a definitive calculation of $D$ for this situation, BH01 proceeded by estimating $D$ for in a C/O plasma by the self-diffusion coefficient, $D_s$, of the classical one-component plasma (OCP)."
+ With this. BUOL then estimated the power released. Ly. by sinkine for a fixed profile.," With this, BH01 then estimated the power released, $L_g$, by sinking for a fixed profile."
+ This calculation suggested that sedimentation nieht release sufficient cucreyto impact WD cooling., This calculation suggested that sedimentation might release sufficient energyto impact WD cooling.
+ We now take the next step in addressing this question bv performing a selfconsisteut evolution of both the density. poy. and the WD core temperature. ων in WDs composed of a single dominant ion species.," We now take the next step in addressing this question by performing a self-consistent evolution of both the density, $\rho_{22}$, and the WD core temperature, $T_c$, in WDs composed of a single dominant ion species."
+ We fiud that heating delays the time it takes a WD to cool to a given lhuuinosity. imass., We find that heating delays the time it takes a WD to cool to a given luminosity. .
+. We also investigate the scisimological impact of the, We also investigate the seismological impact of the
+as or>0.,as $r\rightarrow0$.
+ Note that solutions failing to satisfy this condition sometimes are also acceptable., Note that solutions failing to satisfy this condition sometimes are also acceptable.
+ For example. when the right-hand side of the above equation approaches a finite coustaut. the sineularity at r=0 may. be related to apoint-like particle [|18|..," For example, when the right-hand side of the above equation approaches a finite constant, the singularity at $r=0$ may be related to apoint-like particle \cite{VS}."
+ However. since here we are mainly interested in eravitational collapse. in this paper we shall assume that this coudition holds strictly at the beginning of the collapse. so that we can be sure that the singularity to be formed later on the axis is due to the collapse. (," However, since here we are mainly interested in gravitational collapse, in this paper we shall assume that this condition holds strictly at the beginning of the collapse, so that we can be sure that the singularity to be formed later on the axis is due to the collapse. ("
+idi) No closed. timelike curves.,iii) No closed timelike curves.
+ In spacetimes with circular sviunetry. closed timelike curves cau be easily introduced.," In spacetimes with circular symmetry, closed timelike curves can be easily introduced."
+ To eusure their absence. we assume that the condition holds in the whole spacetime.," To ensure their absence, we assume that the condition holds in the whole spacetime."
+ Iu addition to these conditions. it is usually also required that the spacetime be asviuptotically flat in the radial direction.," In addition to these conditions, it is usually also required that the spacetime be asymptotically flat in the radial direction."
+ However. since we consider solutions with selfsimilarityv. this condition cannot be satistied. unless we restrict thei validity only up to a maximal radius. saver =rg(f). aud then join them with others in the region rzry(t). which are asviuptoticallv fat as r+x.," However, since we consider solutions with self-similarity, this condition cannot be satisfied, unless we restrict their validity only up to a maximal radius, say, $r=r_0(t)$, and then join them with others in the region $r>r_0(t)$, which are asymptotically flat as $r\rightarrow\infty$."
+" Iu this paper. we shall uot consider such a possibility. aud simply assiuue that the sclfsimilay solutions are valid in the whole spacetiuc,"," In this paper, we shall not consider such a possibility, and simply assume that the self-similar solutions are valid in the whole spacetime."
+ Iu this section. we study solutions with sclfsimuilarity of the first kind.," In this section, we study solutions with self-similarity of the first kind."
+" To obtain the desired equatious we substitute equations (3)). CL)) and (5)) iuto the Einstein field equatious and also using the Klein-Crordon equatiou where # is the Eiusteiu coupling coustaut. L1=(/7//7/7VW,V with V, being the covariaut derivative."," To obtain the desired equations we substitute equations \ref{Tmn1}) ), \ref{Tmn2}) ) and \ref{ug}) ) into the Einstein field equations and also using the Klein-Gordon equation where $\kappa$ is the Einstein coupling constant, $\dal=g^{\alpha\beta}\nabla_\alpha\nabla_\beta$ with $\nabla_\alpha$ being the covariant derivative."
+" When a=1. according to equation (10)) terms with powers of kr can be substituted by the same powers of (tf) since r=e""(f)."," When $\alpha=1$, according to equation \ref{x}) ) terms with powers of $r$ can be substituted by the same powers of $(-t)$ since $r=\ex^x(-t)$."
+" Then it can be shown that the Einstein feld equations in this case become (see Appendix A) 4 ll P. 5(00,10,E]] νι. PI) py =0(", Then it can be shown that the Einstein field equations in this case become (see Appendix A) ] ] _0 ] p_0 ] p_0 ] = 0.
+21) Iu order to deteriunuate the metric cocficicuts aud the .& dependence of the scalar field) completely. and to make possible the calculations. we assiunoe the equation of state for the perfect fluid eiven by," In order to determinate the metric coefficients and the $x$ dependence of the scalar field completely, and to make possible the calculations, we assume the equation of state for the perfect fluid given by ,"
+"relativistic beaming effect for ultra relativistic expansion, the observer only sees light travelling towards him, i.e. in the direction (0=¢,y1/2).","relativistic beaming effect for ultra relativistic expansion, the observer only sees light travelling towards him, i.e. in the direction $(\vartheta=\zeta, \varphi=\pi/2)$."
+" For particles located in the current sheets, this corresponds to an observer ts satisfying where {€Z is a natural integer."," For particles located in the current sheets, this corresponds to an observer $~t_{\rm s}$ satisfying where $l\in\mathbb{Z}$ is a natural integer."
+" Solving for this time t£; at which the photon leaves the current sheet, we find The + sign labels the arm of the double spiral structure (seen in the equatorial plane) responsible for the radiation."," Solving for this time $t_s$ at which the photon leaves the current sheet, we find The $\pm$ sign labels the arm of the double spiral structure (seen in the equatorial plane) responsible for the radiation."
+ Recall that the double peak light-curves follow from this double spiral shape., Recall that the double peak light-curves follow from this double spiral shape.
+ The time separation between two consecutive pulses becomes therefore Note that this is not necessarily the high-energy peak separation because it can in principle vary between zero and a full period P., The time separation between two consecutive pulses becomes therefore Note that this is not necessarily the high-energy peak separation because it can in principle vary between zero and a full period $P$.
+" To be more precise, we have to compare one specified pulse with its two neighbors in time, the one coming earlier to the observer and the other later."," To be more precise, we have to compare one specified pulse with its two neighbors in time, the one coming earlier to the observer and the other later."
+" Doing so we get Normalizing to the period of the pulsar P, the phase separation reads From the definition of the arccos function, we get Finally, the relation between obliquity x, line of sight inclination angle C and peak separation A can be rearranged into To summarize we find This relation is shown in Fig."," Doing so we get Normalizing to the period of the pulsar $P$, the phase separation reads From the definition of the arccos function, we get Finally, the relation between obliquity $\chi$, line of sight inclination angle $\zeta$ and peak separation $\Delta$ can be rearranged into To summarize we find This relation is shown in Fig."
+ 1 where cos¢ is plotted versus cosxfor constant value of the separation A starting from zero to 0.5 with a step 0.02.," \ref{fig:SeparationPic} where $\cos
+\zeta$ is plotted versus $\cos \chi$for constant value of the separation $\Delta$ starting from zero to $0.5$ with a step $0.02$."
+ A few special cases are worth to mention., A few special cases are worth to mention.
+" First, zero separation or overlapping of both pulses, corresponding to A=0, implies ¢=7/2—x."," First, zero separation or overlapping of both pulses, corresponding to $\Delta=0$, implies $\zeta
+= \pi/2-\chi$."
+" For the chosen variables in Fig. 1,"," For the chosen variables in Fig. \ref{fig:SeparationPic},"
+ it corresponds to the quarter circle of radius unity centered at the origin., it corresponds to the quarter circle of radius unity centered at the origin.
+" Second, a separation of half a period A=0.5, implies a line of sight contained in the equatorial plane of the pulsar, ¢=7/2."," Second, a separation of half a period $\Delta=0.5$, implies a line of sight contained in the equatorial plane of the pulsar, $\zeta = \pi/2$."
+ It is independent of the obliquity., It is independent of the obliquity.
+ This corresponds to the straight horizontal line along the x-axis in Fig. 1.., This corresponds to the straight horizontal line along the x-axis in Fig. \ref{fig:SeparationPic}. .
+ The relation Eq. (24)), The relation Eq. \ref{eq:SeparationPic6}) )
+" was already noticed by Kirk (2005),, see figure therein."," was already noticed by \cite{2005MmSAI..76..494K}, see figure therein."
+" For symmetry reasons, we will assume that x€[0,7/2], therefore cotx>0."," For symmetry reasons, we will assume that $\chi \in [0,\pi/2]$, therefore $\cot\chi\ge0$."
+" Next, we show that the combination polar cap/striped wind allows to derive a simple analytical relation between the lag of radio vs gamma-ray arrival time ó and the high-energy peak separation A."," Next, we show that the combination polar cap/striped wind allows to derive a simple analytical relation between the lag of radio vs gamma-ray arrival time $\delta$ and the high-energy peak separation $\Delta$."
+" Because the two sites of emission are very distinct, inside and outside the light cylinder, the lag is interpreted as a time of flight delay between the region of radio radiation and the one for gamma-ray photon production."," Because the two sites of emission are very distinct, inside and outside the light cylinder, the lag is interpreted as a time of flight delay between the region of radio radiation and the one for gamma-ray photon production."
+" The choice of origin of time and line of sight located in the plane (Oyz) (therefore — 5) implies that the radio photon emitted towards the observer will happen at periodic dates with period P, for one pole, given by where k€Z is a natural integer."," The choice of origin of time and line of sight located in the plane $(Oyz)$ (therefore $\varphi = \frac{\pi}{2}$ ) implies that the radio photon emitted towards the observer will happen at periodic dates with period $P$, for one pole, given by where $k\in\mathbb{Z}$ is a natural integer."
+" The superscript n denotes bundles of photons emanating from the same polar cap, let us say the north magnetic pole."," The superscript n denotes bundles of photons emanating from the same polar cap, let us say the north magnetic pole."
+" The same applies to the south cap but due to the symmetry, the phase difference corresponds to half a period, P/2, thus 'The observer remains at adistance D from the center of the neutron star."," The same applies to the south cap but due to the symmetry, the phase difference corresponds to half a period, $P/2$ , thus The observer remains at adistance $D$ from the center of the neutron star."
+ The polar cap photon needs therefore a, The polar cap photon needs therefore a
+Maronctal.(2012.hereafterPaperI) (Stoneetal.2008).. Toth(2000).. Ryu&Jones(1995)..,"\citet[][hereafter Paper I]{phurbasalg} \citep{2008ApJS..178..137S}, \cite{2000JCoPh.161..605T}, \cite{1995ApJ...442..228R}."
+ 77 (Springel2005) 3.. LL. (Springel2005)..," \ref{sec_implementation} \citep{2005MNRAS.364.1105S} \ref{sec_tests}. \ref{sec_discussion}. \citep{2005MNRAS.364.1105S},"
+ module described iu Paper I has been incorporated mto a inodified version of that has been re-purposed to serve as a parallel framework., module described in Paper I has been incorporated into a modified version of that has been re-purposed to serve as a parallel framework.
+ We rofer to the version documented iun this work as Phurbas version 1.1 to allow for the differcutiation of future iiocdificatious to the algorithin aud the code. (, We refer to the version documented in this work as Phurbas version 1.1 to allow for the differentiation of future modifications to the algorithm and the code. (
+We note that tests of Pliurbas 1.0 were described in the first subinitted version of this paper. posted to ArXiv as version 1 of this paper.),"We note that tests of Phurbas 1.0 were described in the first submitted version of this paper, posted to ArXiv as version 1 of this paper.)"
+ We sunuuarize here the algorithia described iu detail in Paper 1. Phurbas solves the equations of ideal ATID. expressed im terms of Lagrangian tine derivatives.," We summarize here the algorithm described in detail in Paper I. Phurbas solves the equations of ideal MHD, expressed in terms of Lagrangian time derivatives."
+ The equations are discretized on au adaptive set of particles that carry values of the feld variables (clensity p. velocity V. maguetic field B. aud internal enerew density 0).," The equations are discretized on an adaptive set of particles that carry values of the field variables (density $\rho$, velocity $\sV$, magnetic field $\sB$, and internal energy density $\sigma$ )."
+ The particles move with the velocity V that thev cary. making the Laeranegian formali the natural description of the time evolution of the field variables.," The particles move with the velocity $\sV$ that they carry, making the Lagrangian formalism the natural description of the time evolution of the field variables."
+ The evolution equations relate time derivatives of the field variables to their spatial derivatives., The evolution equations relate time derivatives of the field variables to their spatial derivatives.
+ To calculate the spatial derivatives. Plurbas uses a local. third-order. polvnouual. moving least squares interpolation. using ucighbor particle values drawn from a sphere of radius re=2.3A; about particle 7. where A; is the effective resolution paracter.," To calculate the spatial derivatives, Phurbas uses a local, third-order, polynomial, moving least squares interpolation, using neighbor particle values drawn from a sphere of radius $r_{f,i}=2.3\lambda_i$ about particle $i$, where $\lambda_i$ is the effective resolution parameter."
+ A second order predictor-corrector scheme is used with the time derivatives obtained by applving the MITD equations to the approximate spatial derivatives to advance the particle positions and variable values., A second order predictor-corrector scheme is used with the time derivatives obtained by applying the MHD equations to the approximate spatial derivatives to advance the particle positions and variable values.
+ Particles are added aud deleted. fille voids and destroviug clumps. to ensure that each sphere of radius rg is well sampled.," Particles are added and deleted, filling voids and destroying clumps, to ensure that each sphere of radius $r_f$ is well sampled."
+ On average. the particle creation aud deletion results iu at least one particle per volume A.," On average, the particle creation and deletion results in at least one particle per volume $\lambda^3$."
+ The resolution parameter A need not be coustaut iu space or time. and cach particle has au individual time step.," The resolution parameter $\lambda$ need not be constant in space or time, and each particle has an individual time step."
+ Iu this sense Plhiurbas is fully adaptive both spatially aud temporally., In this sense Phurbas is fully adaptive both spatially and temporally.
+ The time steps are independent of the bulk velocity of the flow., The time steps are independent of the bulk velocity of the flow.
+ Spatial variation of the time steps islunitedto prevent thepenetration of short time step, Spatial variation of the time steps islimitedto prevent thepenetration of short time step
+the extent that they clearly demonstrate a luminosity-dependent clustering trend.,the extent that they clearly demonstrate a luminosity-dependent clustering trend.
+" For now, we turn to independent constraints on optical duty cycles derived from the observed space density of quasars and models of the underlying black hole population."," For now, we turn to independent constraints on optical duty cycles derived from the observed space density of quasars and models of the underlying black hole population."
+ A widely used method to model the accretion history of the BH population employs a continuity equation (Cavaliere et 11972; Small Blandford 1992) to track the growth of the BH mass function that is implied by the observed quasar luminosity function., A widely used method to model the accretion history of the BH population employs a continuity equation (Cavaliere et 1972; Small Blandford 1992) to track the growth of the BH mass function that is implied by the observed quasar luminosity function.
+" This approach is reviewed extensively by Shankar et ((2009a; hereafter SWM), who apply it to a compilation of recent data sets, and whose results and methodology we adopt here."," This approach is reviewed extensively by Shankar et (2009a; hereafter SWM), who apply it to a compilation of recent data sets, and whose results and methodology we adopt here."
+" The parameters of a model are the radiative efficiency e, which converts an observed luminosity to a corresponding mass accretion rate, and the Eddington ratio A=L/Lzaa, which determines the mass of the BHs to be associated with a given observed luminosity."," The parameters of a model are the radiative efficiency $\epsilon$, which converts an observed luminosity to a corresponding mass accretion rate, and the Eddington ratio $\lambda=L/L_{\rm Edd}$, which determines the mass of the BHs to be associated with a given observed luminosity."
+ The method can be generalized to allow a distribution of A values (Shankar 2009)., The method can be generalized to allow a distribution of $\lambda$ values (Shankar 2009).
+" For a single A value, the duty cycle is simply where ®(L,z) is the quasar luminosity function and ©®(Mgu,z) is the BH mass function at the mass that corresponds to luminosity L, Mss=10°\~1(L/10*°! ergs~')Mo."," For a single $\lambda$ value, the duty cycle is simply where $\Phi(L,z)$ is the quasar luminosity function and $\Phi(M_{\rm BH},z)$ is the BH mass function at the mass that corresponds to luminosity $L$, $M_{\rm BH} = 10^8 \lambda^{-1} (L/10^{46.1} \ergsec) M_\odot$ ."
+" For a distribution of A, one must take some care in defining the meaning of the term “active.”"," For a distribution of $\lambda$, one must take some care in defining the meaning of the term “active.”"
+" At redshifts z>1, BH mergers are expected to play a minor role in shaping ®(Mgpu,2) relative to accretion (SWM), and we neglect them here."," At redshifts $z>1$, BH mergers are expected to play a minor role in shaping $\Phi(M_{\rm BH},z)$ relative to accretion (SWM), and we neglect them here."
+" The left panel of Figure 4,, analogous to figure 7c of SWM, shows the optical duty cycle as a function of black hole mass at z=1.45 predicted by several different continuity equation models."," The left panel of Figure \ref{fig|fcont}, analogous to figure 7c of SWM, shows the optical duty cycle as a function of black hole mass at $z=1.45$ predicted by several different continuity equation models."
+" The model shown by the solid curve has A=0.6 and e=0.065, independent of mass and redshift, which SWM show yields a good match to observational estimates of the local black hole mass function."," The model shown by the solid curve has $\lambda=0.6$ and $\epsilon=0.065$, independent of mass and redshift, which SWM show yields a good match to observational estimates of the local black hole mass function."
+" We convert the total duty cycle to the optical duty cycle using fopt=f/3, based on the ratio of the optical luminosity function for the SDSS quasar sample to the bolometric luminosity function in SWM."," We convert the total duty cycle to the optical duty cycle using $\mfopt=f/3$, based on the ratio of the optical luminosity function for the SDSS quasar sample to the bolometric luminosity function in SWM."
+" Above Mau=1053Mo, the predicted duty cycle is fog;=3.6x1073, while the differing shapes of the quasar luminosity function and the evolved black hole mass function imply higher duty cycles at lower masses."," Above $\mmbh=10^{8.8}M_\odot$, the predicted duty cycle is $\mfopt=3.6\times 10^{-3}$, while the differing shapes of the quasar luminosity function and the evolved black hole mass function imply higher duty cycles at lower masses."
+" The thick vertical lines mark the masses that correspond to the lower luminosity limit of the S09 sample at z=1.45, logMpu=8.2—A."," The thick vertical lines mark the masses that correspond to the lower luminosity limit of the S09 sample at $z=1.45$, $\log\mmbh = 8.2-\log\lambda$."
+" As discussed by SWM, including black hole mergers in the mass function evolution or varying the bolometric luminosity functions or bolometric corrections within observationally acceptable bounds has minimal impact on the inferred duty cycles at these redshifts; the largest systematic uncertainties are associated with the choices of A and e."," As discussed by SWM, including black hole mergers in the mass function evolution or varying the bolometric luminosity functions or bolometric corrections within observationally acceptable bounds has minimal impact on the inferred duty cycles at these redshifts; the largest systematic uncertainties are associated with the choices of $\lambda$ and $\epsilon$."
+" The dotted curve shows a model with A=0.3, which has similar shape but higher normalization."," The dotted curve shows a model with $\lambda=0.3$, which has similar shape but higher normalization."
+" The normalization trend is easily understood: the integrated quasar emissivity determines the total mass density of the black hole population (Soltaan 1982), and assuming lower A shifts this density to more massive, hence rarer, black holes, which must have a higher duty cycle to reproduce the luminosity function."," The normalization trend is easily understood: the integrated quasar emissivity determines the total mass density of the black hole population (Sołtaan 1982), and assuming lower $\lambda$ shifts this density to more massive, hence rarer, black holes, which must have a higher duty cycle to reproduce the luminosity function."
+" The dashed curve shows a model with a spread in Eddington ratios, Gaussian in logA with 0.6-dex dispersion and peak at Amed=0.3, evolved with the techniques described in Shankar (2009)."," The dashed curve shows a model with a in Eddington ratios, Gaussian in $\log\lambda$ with 0.6-dex dispersion and peak at $\lambda_{\rm med}=0.3$, evolved with the techniques described in Shankar (2009)."
+ Results are intermediate between the two constant-A models., Results are intermediate between the two $\lambda$ models.
+" However, in this case the S09 luminosity threshold does not correspond to a sharp mass cut, so the dot-dashed curveshows fopt with the additional criterion"," However, in this case the S09 luminosity threshold does not correspond to a sharp mass cut, so the dot-dashed curveshows $\mfopt$ with the additional criterion"
+To motivate our study. we contrast planet formation around low mass stars and stars.,"To motivate our study, we contrast planet formation around low mass stars and solar-type stars."
+ For solar-tvpe stus approaching the main sequence. the luminosity is roughly constant on tvpical planet formation timescales of MMyr.," For solar-type stars approaching the main sequence, the luminosity is roughly constant on typical planet formation timescales of Myr."
+ Thus. the conditions where planets form change little with time.," Thus, the conditions where planets form change little with time."
+ For stars with masses <0.5M... however. the Iuminosity [ades bv a factor of 10.100 on the ILavashi track.," For stars with masses $\lesssim 0.5\,M_{\odot}$, however, the luminosity fades by a factor of 10–100 on the Hayashi track."
+ Because the inner disk radius is locked a a fixed distance relative to the radius of the central star. the inner disk contracts as the star contracts.," Because the inner disk radius is `locked' at a fixed distance relative to the radius of the central star, the inner disk contracts as the star contracts."
+ During this evolution. (he ‘snow line the point that separates the inner region of rocky planet Iormation from the outer region of icv planet formation also moves inward.," During this evolution, the `snow line' – the point that separates the inner region of rocky planet formation from the outer region of icy planet formation – also moves inward."
+ Coupled with the evolution of the inner disk. (he moving snow line produces a dramatic variation in the surface density at fixed distances from the central star.," Coupled with the evolution of the inner disk, the moving snow line produces a dramatic variation in the surface density at fixed distances from the central star."
+ This behavior enables the rapid formation of 1ον protoplanets., This behavior enables the rapid formation of icy protoplanets.
+ As the low mass star approaches the main sequence. these protoplanets collide and merge into super-Earths with properües similar to those detected in recent microlensing experiments.," As the low mass star approaches the main sequence, these protoplanets collide and merge into super-Earths with properties similar to those detected in recent microlensing experiments."
+ In coagulation models. planets grow Irom repeated collisions and mergers of small objects in a circumstellar disk (Safronoy1969)..," In coagulation models, planets grow from repeated collisions and mergers of small objects in a circumstellar disk \citep{1969QB981.S26......}."
+ When 110 km ‘planetesimals’ form and start to erow (Weidenschiling1980:Dullemond&Dominik2005).. dvnamical friction damps the orbital eccentricities of the largest objects.," When 1–10 km `planetesimals' form and start to grow \citep{1980Icar...44..172W,2005A&A...434..971D}, dynamical friction damps the orbital eccentricities of the largest objects."
+ Damping vields large gravitational cross-sections and leads to runaway growth. where the largest objects grow fastest ancl run away [rom more slowly growing smaller objects (Wetherill&Stewart1989:IXokuboIda1996).," Damping yields large gravitational cross-sections and leads to `runaway growth,' where the largest objects grow fastest and run away from more slowly growing smaller objects \citep{1989Icar...77..330W,1996Icar..123..180K}."
+. Throughowt the runaway. the largest protoplanets stir up the leftover planetesimals.," Throughout the runaway, the largest protoplanets stir up the leftover planetesimals."
+ Eventually. the leftovers have orbital velocity dispersions comparable to the escape velocities of the largest protoplanets.," Eventually, the leftovers have orbital velocity dispersions comparable to the escape velocities of the largest protoplanets."
+ Because gravitational eross-sections fall as velocity dispersions rise. runaway erowth ends.," Because gravitational cross-sections fall as velocity dispersions rise, runaway growth ends."
+ The ensemble of planetesimals aud protoplanets (hen enters 'oligarchic growth. where the largest objects oligarchs αστείο alrates roughly independent of their size (Ixokubo&Ida1993)..," The ensemble of planetesimals and protoplanets then enters `oligarchic' growth, where the largest objects – oligarchs – accrete atrates roughly independent of their size \citep{1998Icar..131..171K}."
+ Duringe oligarchice exgrowth. protoplanets become isolated [rom their surroundings.," During oligarchic growth, protoplanets become isolated from their surroundings."
+"e If an oligarch accretes all of the mass in an annulus with width 25 Hj. where Ry=«(M/3M,jy "," If an oligarch accretes all of the mass in an annulus with width $2 B R_H$ , where $R_H = a (M / 3 M_{\star})^{1/3}$ "
+the binary. ancl the detailed evolutionary phases of the svstem prior to SN explosion are currently the matter of strong debate (see. lor instance. Canal. Ménndez. Iuiz-Lapuente 2001. and references therein).,"the binary, and the detailed evolutionary phases of the system prior to SN explosion are currently the matter of strong debate (see, for instance, Canal, Ménndez, Ruiz-Lapuente 2001, and references therein)."
+ Those uncertainties raise doubts about the use of high-redshift σα for cosmological purposes. especially as their calibration is only based on the properties of local SNla.," Those uncertainties raise doubts about the use of high-redshift SNIa for cosmological purposes, especially as their calibration is only based on the properties of local SNIa."
+ One wax to discern among the various presupernova models is to look for their consequences on supernova remnant evolution., One way to discern among the various presupernova models is to look for their consequences on supernova remnant evolution.
+ Models of supernovae remnants (SNR) originated bv à SNla usually assume either a uniform density interstellar medium (ISM) or a pxr7 cireunstellar medium (CSM). produced by a constant mass loss rate wind (Dwarkadas&Chevalier1998.[orinstance)..," Models of supernovae remnants (SNR) originated by a SNIa usually assume either a uniform density interstellar medium (ISM) or a $\rho\propto r^{-2}$ circumstellar medium (CSM), produced by a constant mass loss rate wind \citep[for instance]{dc98}."
+ This assumption. together wilh an appropriate choice of the SN ejecta profile. allows Lor the existence of similarity solutions that describe the early evolution of the SNR (Chevalier&Mcelxee 1999).," This assumption, together with an appropriate choice of the SN ejecta profile, allows for the existence of similarity solutions that describe the early evolution of the SNR \citep{ch82,tm99}."
+. Detailed models of pre-SNla binary svstem evolution. however. point to a lime varying mass loss rate (ILachisu.Kato.al. 2000).. which is expected to lead to ambient medium (AM) density profiles very clilferent from a power las.," Detailed models of pre-SNIa binary system evolution, however, point to a time varying mass loss rate \citep{hkn96,lan00}, which is expected to lead to ambient medium (AM) density profiles very different from a power law."
+ Verv lew observations have been made during the initial phase of formation οἱ supernovae remnants. when the interaction of supernova ejecta wilh presupernova wind could be tested.," Very few observations have been made during the initial phase of formation of supernovae remnants, when the interaction of supernova ejecta with presupernova wind could be tested."
+ Schlegel&Petre(1993) used ROSAT to search for X-ray emission from the (vpe la SNL1992A 16 days alter visual maximum. ancl derived an upper limit for the presupernova mass-loss rate of less than a few times LOM.vr.|. Cumm," \citet{sp93} used ROSAT to search for X-ray emission from the type Ia SN1992A 16 days after visual maximum, and derived an upper limit for the presupernova mass-loss rate of less than a few times $10^{-6}
+M_{\odot}{\rm yr}^{-1}$."
+ingetal.(1996) looked for narrow Io in a high-resolution spectrum of $N1994D. another SNIa. and derived an upper limit for the mass loss rate of ~1.5xLOM.vr.|.," \citet{cum96} looked for narrow $\alpha$ in a high-resolution spectrum of SN1994D, another SNIa, and derived an upper limit for the mass loss rate of $\sim1.5\times10^{-5} M_{\odot}{\rm yr}^{-1}$."
+ These observations have in common that they were made αἱ a verv early phase in the supernova ejecta evolution. implving that thev only probed the CSM very close to the progenitor.," These observations have in common that they were made at a very early phase in the supernova ejecta evolution, implying that they only probed the CSM very close to the progenitor."
+ The AM region up lo a lew parsecs can be probed by the evolution of voung SNR., The AM region up to a few parsecs can be probed by the evolution of young SNR.
+atmospheres. both the effects of scattering and line dauketiug contribute to ¢ Wagoner1981:Eastiuauctal. 1996).,"atmospheres, both the effects of scattering and line blanketing contribute to $\zeta$ \citep{Wagoner_EPM,Eastman_EPM}. ."
+. To deteiiuine £ using the expanding photosphere nethod. the observer iieasires aand the time since explosion phf. and estimates the ohotospherie temperature from the color of the spectram.," To determine $L$ using the expanding photosphere method, the observer measures and the time since explosion $t$, and estimates the photospheric temperature from the color of the spectrum."
+"Tj, The dilution factor must )o calculated using detailed uuncerical models (the main complexity of the approach).", The dilution factor must be calculated using detailed numerical models (the main complexity of the approach).
+ NLTE spectral modeling finds that ¢ varies hetween 0.5 aud 2.0. aud is chieflv a uction of huninosity. being rather insensitive to other ejecta paraicters such as the density structure (Eastinauetal.1996:Dessart&Uallicr 2005a).," NLTE spectral modeling finds that $\zeta$ varies between 0.5 and 2.0, and is chiefly a function of luminosity, being rather insensitive to other ejecta parameters such as the density structure \citep{Eastman_EPM, Dessart_EPM}."
+. The standard candle relation is simply au expression of Eq., The standard candle relation is simply an expression of Eq.
+ 16 under certain restricted conditions., \ref{Eq:Lstef} under certain restricted conditions.
+ The time since explosion f is. by coustruction. fixed at 50. days.," The time since explosion $t$ is, by construction, fixed at 50 days."
+ The temperature for SNe TIP ou Tythe plateau is uecarly a constant. constrained to be near the recombination temperature T;6000 Ix. The dilution factor ¢ may vary from eveut to event. but if¢ is primarily a function of huuinositv this dependence can be absorbed into the exponent.," The temperature for SNe IIP on the plateau is nearly a constant, constrained to be near the recombination temperature $T_i \approx 6000$ K. The dilution factor $\zeta$ may vary from event to event, but if $\zeta$ is primarily a function of luminosity this dependence can be absorbed into the exponent."
+" This nuplics £L=Cr, where the coustaut C and the non-blackbody effects € can be calibrated using a sample of nearby objects. or a set of theoretical models."," This implies $L = C \vph^{2+\epsilon}$, where the constant $C$ and the non-blackbody effects $\epsilon$ can be calibrated using a sample of nearby objects, or a set of theoretical models."
+ The SC relation need not be applied only at day 50. and we fiud that similar roelatious apply all alone the plateau.," The SC relation need not be applied only at day 50, and we find that similar relations apply all along the plateau."
+ However the time since explosion iust be known as the jorinalizatiou depends ou time (Eq. 16))., However the time since explosion must be known as the normalization depends on time (Eq. \ref{Eq:Lstef}) ).
+ We fud that an uncertaiutv in explosion time of LO davs leads to an error iu inferred brightness of 0.2.—0.3 mae., We find that an uncertainty in explosion time of 10 days leads to an error in inferred brightness of $0.2-0.3$ mag.
+" It is unise o apply the SC relation at times much earlier than 30 davs. as the ejecta temperatures are likely too high for reconibiation to have set iu. aud there is no assurance hat Ty,71."," It is unwise to apply the SC relation at times much earlier than 30 days, as the ejecta temperatures are likely too high for recombination to have set in, and there is no assurance that $\Tph \approx T_i$."
+ One nice feature of the models is that thev offer an absolute normalization of the SC relation without weeding to assume a value of the ITubble coustaut., One nice feature of the models is that they offer an absolute normalization of the SC relation without needing to assume a value of the Hubble constant.
+ By fitting the relation evaluated at different times since explosion. we fiud The models do predicta deviation from the simple LxCh relation of Eq. 16..," By fitting the relation evaluated at different times since explosion, we find The models do predicta deviation from the simple $L \propto \vph^2$ relation of Eq. \ref{Eq:Lstef},"
+ showing instead £Xc general accordance with that found in the observational saluple (Παπ&Pinto2002)..," showing instead $L \propto
+\vph^{2.75}$ in general accordance with that found in the observational sample \citep{Hamuy_SCM}."
+ This effect is primarily due to the deviation of the spectrum from a blackbody., This effect is primarily due to the deviation of the spectrum from a blackbody.
+ The aodel relation of Figure 165 las a simular normalization to the observations. taken from ILhuuuxy(2003)..," The model relation of Figure \ref{Fig:Hamuy} has a similar normalization to the observations, taken from \cite{Hamuy_SNIIP}."
+ This implies that our model SC relation is iu rough agreement with the distauces to SNe TIP obtained in other wavs., This implies that our model SC relation is in rough agreement with the distances to SNe IIP obtained in other ways.
+ Particularly comforting is theagreement with SN 19990. which has a measured Cepheid distance to its host ealaxy NGC 1637 of 11.741.0 Mpe (Leonardetal.2003).," Particularly comforting is theagreement with SN 1999em, which has a measured Cepheid distance to its host galaxy NGC 1637 of $11.7 \pm 1.0$ Mpc \citep{Leonard_ceph}."
+. We find a very simular distance of 11.6+4 Alpe from Eq., We find a very similar distance of $11.6 \pm 1.2$ Mpc from Eq.
+" 17. wheu taking the observed values my,=L398. μι=0757kus and (followingBaronetal.2000:ILEuuuyv 2001).. an extinction of A,=(0.31."," \ref{Eq:SCR} when taking the observed values $m_v = 13.98$, $\vph = 3757~\kms$, and \citep[following][]
+{Baron_2000, Hamuy_EPM}, an extinction of $A_v = 0.31$."
+ This distance is also consistent witli independent estimates using the expanding plotospher« ucthod (Dessart&Illlier2005a)) aud SEAAL (Baronetal.2001).., This distance is also consistent with independent estimates using the expanding photosphere method \citep{Dessart_EPM} and SEAM \citep{Baron_99em}.
+ One drawback of the standard caudle method. from he observational point of view. is that a high quality spectra is needed to measure the photospheric velocity a difficult prospect for high redshift eveuts.," One drawback of the standard candle method, from the observational point of view, is that a high quality spectrum is needed to measure the photospheric velocity – a difficult prospect for high redshift events."
+ As future survevs will observe light curves for a enormous uunmber of SNe IIP with limited spectroscopic follow-up. methods of purely photometric calibration. however coarse. may be of interest.," As future surveys will observe light curves for a enormous number of SNe IIP with limited spectroscopic follow-up, methods of purely photometric calibration, however coarse, may be of interest."
+ As the explosion cucreyv is the primary variable determine both the plateau liminosity aud duration. we explored the relationship between these two observables.," As the explosion energy is the primary variable determining both the plateau luminosity and duration, we explored the relationship between these two observables."
+ A velatiouship exists (Figure 17)) and is fit by Applying this relation reduces the dispersion from linag down to 0. Linag., A relationship exists (Figure \ref{Fig:ltp}) ) and is fit by Applying this relation reduces the dispersion from 1 mag down to 0.4 mag.
+" Iu practice. the measured plateau duration f, mast be corrected for the effect of the ejected nunass on its duration i order to doteruinue t,,.0."," In practice, the measured plateau duration $t_p$ must be corrected for the effect of the ejected mass on its duration in order to determine $t_{p,0}$."
+ The residual scatter in the relation is clearly due to variation in progenitor mitial mass or metallicity for a eiven explosion energy., The residual scatter in the relation is clearly due to variation in progenitor initial mass or metallicity for a given explosion energy.
+ Prestunably. the scatter could be reduced further by using additional light curve relation. such as the color evolution.," Presumably, the scatter could be reduced further by using additional light curve relation, such as the color evolution."
+ We explored the light curves aud spectra of SNe II models with serious progenitor masses. mctallicitics. and explosion energies.," We explored the light curves and spectra of SNe II models with various progenitor masses, metallicities, and explosion energies."
+" We found that explosions with enereies 0.9.LS D of stars with initial masses in the range 1525AL. can explain the observed rauge of ""uuimosities. velocities. aud ποτ curve duratious of nost SNe IIP."," We found that explosions with energies $0.3-4.8$ B of stars with initial masses in the range $12-25~\Msun$ can explain the observed range of luminosities, velocities, and light curve durations of most SNe IIP."
+ For existing aud future observational suvevs. he model results should be useful forinferrime the xoesenitor star properties. explosion energies. distances. and dust extinction of observed events.," For existing and future observational surveys, the model results should be useful forinferring the progenitor star properties, explosion energies, distances, and dust extinction of observed events."
+ This study. as have previous studies. quantified. how he basic supernova parameters Ro. aud E) affect he light curves.," This study, as have previous studies, quantified how the basic supernova parameters , $R_0$ , and $E$ ) affect the light curves."
+ We also highlighted (M.the iniportaut role of two additional parameters: the radioactive, We also highlighted the important role of two additional parameters: the radioactive
+"The energv released in the accretion disk per unit surface can be caleulated by using the standard disk theory and is written as (INato. Fukue. Mineshige 1993) where M. M. and ry, is the mass of black hole. the accretion rate. and the inner boundary raclius of the accretion disks where (he torque vanishes. respectively,","The energy released in the accretion disk per unit surface can be calculated by using the standard disk theory and is written as (Kato, Fukue, Mineshige 1998) where $M$ , $\dot{M}$ , and $r_\mathrm{in}$ is the mass of black hole, the accretion rate, and the inner boundary radius of the accretion disks where the torque vanishes, respectively."
+ We take Af=3.\/. and Mn=35Q~3xLO°em where ry is the Schwarzschild radius., We take $M = 3 M_{\odot}$ and $r_\mathrm{in} = 3 r_\mathrm{s} \simeq 3 \times 10^6 \mathrm{cm}$ where $r_\mathrm{s}$ is the Schwarzschild radius.
+ This dissipated energy is. if the neutrino cooling is ellicient. released as neutrino emission.," This dissipated energy is, if the neutrino cooling is efficient, released as neutrino emission."
+ We consider (he aud homogeneous corona wilh vertical thickness // above the disk (Fig.1))., We consider the plane-parallel and homogeneous corona with vertical thickness $H$ above the disk \ref{fig:corona}) ).
+ With conservalive views (he parameter value // may be comparable to the disk thickness. since ihe most unstable wavelength of Parker instability is nearly scale height in the disk (e.g. Matsumoto et al.," With conservative views the parameter value $H$ may be comparable to the disk thickness, since the most unstable wavelength of Parker instability is nearly scale height in the disk (e.g. Matsumoto et al."
+ L988) and thus. the typical scale of (he magnetic field in (he corona would be comparable to the scale height.," 1988) and thus, the typical scale of the magnetic field in the corona would be comparable to the scale height."
+ Neutrino emission from (he inner part of the disk dominates over that [vom (he outer part and mainly contributes to the heating of relativistic fireball above the disk., Neutrino emission from the inner part of the disk dominates over that from the outer part and mainly contributes to the heating of relativistic fireball above the disk.
+ Then. we evaluate the enhancement of mean οποιον of neutrinos al r=dr. where the energy dissipation rate reaches nearly a maximum value. ancl we assume (hat (his enhancement is proportional to that of the total energy deposition rate by neutrino annihilation.," Then, we evaluate the enhancement of mean energy of neutrinos at $r = 4 r_\mathrm{s}$, where the energy dissipation rate reaches nearly a maximum value, and we assume that this enhancement is proportional to that of the total energy deposition rate by neutrino annihilation."
+ We assume that the corona consists of pure relativistic electron-positron plasma since (he corona nav form above the surface of the disk where the density of barvon is much less than that inside ol the disk (discussed later)., We assume that the corona consists of pure relativistic electron-positron plasma since the corona may form above the surface of the disk where the density of baryon is much less than that inside of the disk (discussed later).
+ We also assume that the electrons and positrons are completely ihermalized. since the timescale of electromagnetic interactions is much shorter than that of weak interactions bv many order of magnitude.," We also assume that the electrons and positrons are completely thermalized, since the timescale of electromagnetic interactions is much shorter than that of weak interactions by many order of magnitude."
+ As the coronal cooling processes we take into account neutrino reactions as The thermal neutrinos emitted from the disk are wp-scatterecd by hot electrons aud positrons in the corona., As the coronal cooling processes we take into account neutrino reactions as The thermal neutrinos emitted from the disk are up-scattered by hot electrons and positrons in the corona.
+ The scattered neutrinos have higher energy since the temperature of the corona (1) is higher than that of the disk (Z4) aud. hence. emerged neutrino spectrum is delormed.," The scattered neutrinos have higher energy since the temperature of the corona $T_\mathrm{c}$ ) is higher than that of the disk $T_\mathrm{d}$ ) and, hence, emerged neutrino spectrum is deformed."
+ The corona is cooled via scatterings of neutrinos by high energy electrons and positrons in (hecorona., The corona is cooled via scatterings of neutrinos by high energy electrons and positrons in thecorona.
+Some part of scattered neutrinos are re-absorbed and heat the disk.,"Some part of scattered neutrinos are re-absorbed and heat the disk,"
+calculate the correlation between the GOLF observations and the two activity proxies.,calculate the correlation between the GOLF observations and the two activity proxies.
+ Figure 5 shows p-mode velocity amplitudes variation as seen by GOLF red-wing configuration over the period 1996-2001 compared to the changes observed in ISN and Mgll., Figure \ref{fig:golf_act} shows $p$ -mode velocity amplitudes variation as seen by GOLF red-wing configuration over the period 1996-2001 compared to the changes observed in ISN and MgII.
+" It clearly shows that the size of the p-mode velocity amplitude suppression increases with increasing level of magnetic activity for both activity proxies in two out of three p-bands,."," It clearly shows that the size of the $p$ -mode velocity amplitude suppression increases with increasing level of magnetic activity for both activity proxies in two out of three $p$ -bands,."
+" We then calculated the correlation coefficient with the Spearman rank formula, because it also provides the corresponding probability of a null correlation."," We then calculated the correlation coefficient with the Spearman rank formula, because it also provides the corresponding probability of a null correlation."
+ Table 3 indicates a correlation of 8096 in two out of the three p-bands for ISN and Mgll activity proxies., Table 3 indicates a correlation of $\%$ in two out of the three $p$ -bands for ISN and MgII activity proxies.
+ These findings agree with previous results (??)..," These findings agree with previous results \citep{Jim04,Cha07}."
+" We used 11 years of the VIRGO data starting from 1996 June 25 until 2006 June 25, while for the GOLF observations we"," We used 11 years of the VIRGO data starting from 1996 June 25 until 2006 June 25, while for the GOLF observations we"
+caused by the presence o ‘low-level exteuded emissiou in the secondary [ux calibrators themselves.,caused by the presence of low-level extended emission in the secondary flux calibrators themselves.
+ The secoud detail to cousider was {1at. by cropping 163347711. we had. no secondary. [Iux density calibrators for the recalibrated seasn 1 data.," The second detail to consider was that, by dropping 1633+741, we had no secondary flux density calibrators for the recalibrated season 1 data."
+ However. our analysis of the light curves ancl the time delay results shiowed that the icluslou of the L633+741 data introduced more scatte‘(hau it took out.," However, our analysis of the light curves and the time delay results showed that the inclusion of the 1633+741 data introduced more scatter than it took out."
+ Thus. we believe that our «lecision to drop the 1633+711 data was the proper ole.," Thus, we believe that our decision to drop the 1633+741 data was the proper one."
+ The final corrected ligit curves Or he four components in B16084-656 are slowh in Figure 2.., The final corrected light curves for the four components in B1608+656 are shown in Figure \ref{fig_lc4comb}.
+ The errors ou the points are a combiiation iu quadrature of (1) the RMS noise it the residual map [rom the processing. tvpicaly —0.1 mJy Lane (9) 10.65€ [ractional in'ertalnty arising [rom errors in the absolute Πιx density calibratio1.," The errors on the points are a combination in quadrature of (1) the RMS noise in the residual map from the processing, typically $\sim$ 0.1 mJy $^{-1}$, and (2) a fractional uncertainty arising from errors in the absolute flux density calibration."
+ We have estimaed the ffractional uncertainty from the season 3 CSO light curves., We have estimated the fractional uncertainty from the season 3 CSO light curves.
+ All of tle ight curves have been oriualizecd by thelr mean flux desities over the entire tluree seasons of oservations., All of the light curves have been normalized by their mean flux densities over the entire three seasons of observations.
+ Thus. tie. plot in Figure represents the [ractional variations in flux «ensily. which alows for clirect. comparison of the component light cuves.," Thus, the plot in Figure \ref{fig_lc4comb} represents the fractional variations in flux density, which allows for direct comparison of the component light curves."
+ The figures show tliat the backerounu ποιο varied siguificantly in both season 2 and seaso 13., The figures show that the background source varied significantly in both season 2 and season 3.
+ The variations in eacl of tl1036. SeasOlls are on the order of .. in contrast to the ~5% varlatlons seen iu seasol ri," The variations in each of these seasons are on the order of , in contrast to the $\sim$ variations seen in season 1."
+le ufortuiately. the season 2 light ¢rves lave nearly a colstalt sope. which mace it clifficul to determine lllaimbiguously the time delays and inagnificatiois ueeded to align the component lielit curves.," Unfortunately, the season 2 light curves have nearly a constant slope, which made it difficult to determine unambiguously the time delays and magnifications needed to align the component light curves."
+ However. the season 3 ight curves have both a large amouit of variation and clear chauges in slope. making tleur excellent inputs [or determiuiug time delays.," However, the season 3 light curves have both a large amount of variation and clear changes in slope, making them excellent inputs for determining time delays."
+ In Paper I we used several statistical methods to determine the delays between the light curves., In Paper I we used several statistical methods to determine the delays between the light curves.
+ In this paper. with obviois source variability aud. ~250 ;»oiuts in the Πο curves of each component. we have used only tle dispersion method described by Peltetal.(1991.1996)..," In this paper, with obvious source variability and $\sim$ 250 points in the light curves of each component, we have used only the dispersion method described by \citet{pelt1,pelt2}."
+" This method has the clear advantage of Lot rec,uring any interpolatio of tje input light curves.", This method has the clear advantage of not requiring any interpolation of the input light curves.
+ As we did in Paper L. we used the D3 ALL Disijethocls. (asclefinedinPeltetal.1996) to compare the three incepenclent pairs of light €wves.," As we did in Paper I, we used the $D^2_2$ and $D^2_{4,2}$ methods \citep[as defined
+in][]{pelt2} to compare the three independent pairs of light curves."
+ Iu hese metlods. one of ie input light curves is shifted iu time by a delay (7). ). scaled iu [lus N areative nagification (4i . alid combined with the second curve to create a Composite Ligh Crye.," In these methods, one of the input light curves is shifted in time by a delay $\tau$ ), scaled in flux by a relative magnification $\mu$ ), and combined with the second curve to create a composite light curve."
+ The iuterial dispersio 1in tle composie curve is then calculated by computing the weighed SUL1 of tle NCuared flux ¢lference between pairs of poit in⋅ the curve., The internal dispersion in the composite curve is then calculated by computing the weighted sum of the squared flux difference between pairs of points in the curve.
+ For4 the D3» iuethod. oiv adjacent yas of points are corsidered. wile. flor the D method. additional pairs of points aiο COLsiclered.," For the $D^2_2$ method, only adjacent pairs of points are considered, while, for the $D^2_{4,2}$ method, additional pairs of points are considered."
+" Li the Di, calculations. all pai* separated fewer than 9 days contribute to the fiial dispersiou. but with au acdditiona weiehtiug such that t pairs with the smallest separation i time have the highest weights."," In the $D^2_{4,2}$ calculations, all pairs separated by fewer than $\delta$ days contribute to the final dispersion, but with an additional weighting such that the pairs with the smallest separation in time have the highest weights."
+" By 1icludius more pairs. t Di, statistic is less likely to be affeced by iudividual discrepaut pots."," By including more pairs, the $D^2_{4,2}$ statistic is less likely to be affected by individual discrepant points."
+We searched a rauge of delays aid relative iagnilicatious to determi1ο the values that. when,"We searched a range of delays and relative magnifications to determine the values that, when"
+detection of SSC emissivity at | keV turns out to be insensitive O uuax fd ies between hese two values. and so we Πχ μις at its larges value.,"detection of SSC emissivity at 1 keV turns out to be insensitive to $\gamma_{\rm max}$ if it lies between these two values, and so we fix $\gamma_{\rm max}$ at its largest value."
+ With these parameters. the equipartition field strengths of he two eastern hotspot components. assuming no contribution to he energy density from non-radiating particles such as protons. are 24 nT (primary) and 16 nT (secondary). and the predicted SSC flux densities at | keV are respectively 0.44 and 2.6 nJy.," With these parameters, the equipartition field strengths of the two eastern hotspot components, assuming no contribution to the energy density from non-radiating particles such as protons, are 24 nT (primary) and 16 nT (secondary), and the predicted SSC flux densities at 1 keV are respectively 0.44 and 2.6 nJy."
+ The predicted johoeton index at this frequency is 1.55 (of course. this is simply a function of the electron energy specrum. and so is true for any inverse-Compton process).," The predicted photon index at this frequency is 1.55 (of course, this is simply a function of the electron energy spectrum, and so is true for any inverse-Compton process)."
+ refflux shows the synchrotron fluxes and SSC prediction for the secondary hotspot., \\ref{flux} shows the synchrotron fluxes and SSC prediction for the secondary hotspot.
+ The predictec SSC flux density for the much weaker western hotspot pair is negligible. ~0.07 nly. corresponding to 2 counts in his observation.," The predicted SSC flux density for the much weaker western hotspot pair is negligible, $\sim 0.07$ nJy, corresponding to 2 counts in this observation."
+" These predictions are relatively insensitive to cosmological parameters: for example. using a cosmology where Quintin=1.0 gives a 3 per cent decrease in the expected flux density from the secondary. whie using Jig= TOkms + Mpe1. 0,)iuller=90.3. Ov=0.7 gives a 2 per cent decrease."," These predictions are relatively insensitive to cosmological parameters; for example, using a cosmology where $\Omega_{\rm matter}
+= 1.0$ gives a 3 per cent decrease in the expected flux density from the secondary, while using $H_0 = 70$ km $^{-1}$ $^{-1}$, $\Omega_{\rm matter} = 0.3$, $\Omega_\Lambda = 0.7$ gives a 2 per cent decrease."
+ The hotspots may be projected. but his also has only a weak effect: if the projection angle is 45° then the true long axis is larger by a factor y/2 and the predicted flux is about 7 per cent lower.," The hotspots may be projected, but this also has only a weak effect: if the projection angle is $45^\circ$ then the true long axis is larger by a factor $\sqrt{2}$ and the predicted flux is about 7 per cent lower."
+ One clear systematic error in the caculations comes from an assumption of spherical symmetry for he scattering geometry which is used in the SSC code., One clear systematic error in the calculations comes from an assumption of spherical symmetry for the scattering geometry which is used in the SSC code.
+ Because a eylinder has a higher surface area to volume ratio. the mean οποίου density is ower in a cylinder than in a sphere for a given volume synchrotron emissivity.," Because a cylinder has a higher surface area to volume ratio, the mean photon density is lower in a cylinder than in a sphere for a given volume synchrotron emissivity."
+ We estimate that this effect is less than [O per cent for the secondary component. but we may be overesimating the SSC flux from the fainter primary by 30 per cent (these values are incre:ised if there is substantial projection).," We estimate that this effect is less than 10 per cent for the secondary component, but we may be overestimating the SSC flux from the fainter primary by 30 per cent (these values are increased if there is substantial projection)."
+ Comparing our predicted flux densities with the data. we see that tye magnitude of he observed flux density. its origin in the larger eastern hotspot and its photon index are all in. good agreement with the equipartition predictions of the SSC model. as is the non-detection of the western hotspot pair.," Comparing our predicted flux densities with the data, we see that the magnitude of the observed flux density, its origin in the larger eastern hotspot and its photon index are all in good agreement with the equipartition predictions of the SSC model, as is the non-detection of the western hotspot pair."
+ The observed keV flux density of the E hotspots is somewjat higher than the predicted value for an equipartition field. though only by about 1.5 standard deviations.," The observed 1-keV flux density of the E hotspots is somewhat higher than the predicted value for an equipartition field, though only by about 1.5 standard deviations."
+ Most of the changes to the model discussed above have the effect of reducing the SSC flux density: to increase it we must reduce the magnetic field strengh or find an additional (external) source of photons., Most of the changes to the model discussed above have the effect of reducing the SSC flux density; to increase it we must reduce the magnetic field strength or find an additional (external) source of photons.
+ If the magnetic field in the secondary hotspot is reduced by 25 per cent to 12 nT. the secondary can produce all the flux seen in X-rays.," If the magnetic field in the secondary hotspot is reduced by 25 per cent to 12 nT, the secondary can produce all the flux seen in X-rays."
+ Neglecting the small SSC contribution from the primary hotspot. the data with their associated uncertainties imply within he SSC model that the magnetic field strength in the secondary hospot is 12óc2 nT (le statistical errors only).," Neglecting the small SSC contribution from the primary hotspot, the data with their associated uncertainties imply within the SSC model that the magnetic field strength in the secondary hotspot is $12 \pm 2$ nT $1\sigma$ statistical errors only)."
+ One possible external source of photons is he active nucleus of 1123., One possible external source of photons is the active nucleus of 123.
+ Inverse-Compton (IC) scattering of photons from the active nucleus will make a signiticant contribuion to the X-ray emission if such photons (at frequencies around LO‘! Hz. since auin7100 )y are comparable in number density o the synchrotron photons.," Inverse-Compton (IC) scattering of photons from the active nucleus will make a significant contribution to the X-ray emission if such photons (at frequencies around $10^{11}$ Hz, since $\gamma_{\rm
+min} \approx 1000$ ) are comparable in number density to the synchrotron photons."
+ At this frequency. the number density of synchrotron photons is approximately 0.08m Hz ὃν since the jotspots are a projected 37 kpe from the nucleus. a similar number density would be produced from the nucleus if its luminosity at tus frequency as seen by the hotspot were <2107* W + !. which would correspond to a flux density of z100 Jy at 1Y Hz. (," At this frequency, the number density of synchrotron photons is approximately 0.08 $^{-3}$ $^{-1}$; since the hotspots are a projected 37 kpc from the nucleus, a similar number density would be produced from the nucleus if its luminosity at this frequency as seen by the hotspot were $\ga 2 \times 10^{27}$ W $^{-1}$ $^{-1}$ which would correspond to a flux density of $\ga
+100$ Jy at $10^{11}$ Hz. ("
+"This condition is equivalent to Si,=(302H)S. «here JJ is the core-hotspot distance and /? is the hotspot radius. and he two fluxes are measured at the required frequency.)","This condition is equivalent to $S_{\rm
+core} = (3D/2R)^2 S_{\rm hs}$, where $D$ is the core-hotspot distance and $R$ is the hotspot radius, and the two fluxes are measured at the required frequency.)"
+ The observed core flux density of 1123 at 1011 Hz is about 40 mJy. though it may be variable (Looney Hardeastle 2000). so that isotropic radiation from the core cannot provide the required photon density.," The observed core flux density of 123 at $10^{11}$ Hz is about 40 mJy, though it may be variable (Looney Hardcastle 2000), so that isotropic radiation from the core cannot provide the required photon density."
+ However. if the core emission at this frequency is beamed. the hotspot will see the core as having a higher luminosity than the one we observe.," However, if the core emission at this frequency is beamed, the hotspot will see the core as having a higher luminosity than the one we observe."
+ We assume the most favourable case for this model of no misalignment between the pe- and kpe-scale jet. though such good alignment is not often observed. and we neglect effects due © the finite angle subtended by the hotspot at the nucleus. which may be significant.," We assume the most favourable case for this model of no misalignment between the pc- and kpc-scale jet, though such good alignment is not often observed, and we neglect effects due to the finite angle subtended by the hotspot at the nucleus, which may be significant."
+" The ratio of required to observed flux. R CR52500 for rough equality of predicted SSC and IC flux densities). hen constrains 2j. the bulk speed in the nucleus. and 8. the angle of the core-hotspot vector to the line of sight: where the term (1—cos8)"" approximately corrects for the anisotropic nature of the resulting IC emission JJones. O'Dell Stein 1974).and the core is treated as a two-sided jet with"," The ratio of required to observed flux, ${\cal R}$ ${\cal R} \approx 2500$ for rough equality of predicted SSC and IC flux densities), then constrains $\beta$, the bulk speed in the nucleus, and $\theta$, the angle of the core-hotspot vector to the line of sight: where the term $(1-\cos \theta)^{-(1+\alpha)}$ approximately corrects for the anisotropic nature of the resulting IC emission Jones, O'Dell Stein 1974),and the core is treated as a two-sided jet with"
+and universality of their distribution are entirely reasonable.,and universality of their distribution are entirely reasonable.
+ Among the three fits to the histogram the best one is that with s=15 km s!., Among the three fits to the histogram the best one is that with $s=15$ km $^{-1}$.
+ Thus. we find that the standard error Avg of a single velocity measurement based on Hy fitting isequal to 1540.5 km s'.," Thus, we find that the standard error $\Delta v_{\rm H}$ of a single velocity measurement based on $_\beta$ fitting isequal to $\pm$ 0.5 km $^{-1}$."
+ A similar accuracy of velocity measurements from VIMOS spectra (10 — 20 km s') was reported by Giuffridaetal.(2010)., A similar accuracy of velocity measurements from VIMOS spectra (10 – 20 km $^{-1}$ ) was reported by \citet{giu10}.
+. Quantitatively the same. but intuitive rather than rigorous estimate of Avy follows from Fig. 5.. ," Quantitatively the same, but intuitive rather than rigorous estimate of $\Delta v_{\rm H}$ follows from Fig. \ref{fig:sigHb_sigNa}, ,"
+in which the histogram of oy peaks at ~15 kms7!., in which the histogram of $\sigma_{\rm H}$ peaks at $\sim$ 15 km $^{-1}$.
+" The stellar Na doublet was usually blended with the interstellar one in such a way that the velocity could be reliably measured only from interstellar D» and stellar D, line (seealsovanLoonetal. 2007).", The stellar Na doublet was usually blended with the interstellar one in such a way that the velocity could be reliably measured only from interstellar $_2$ and stellar $_1$ line \citep[see also][]{vl07}.
+. The interstellar line was measured i1 all 63 objects whose spectra included the Na doublet. yielding the histogram of cw; shown in Fig. 5..," The interstellar line was measured in all 63 objects whose spectra included the Na doublet, yielding the histogram of $\sigma_{\rm Na,i}$ shown in Fig. \ref{fig:sigHb_sigNa}."
+ By analogy with the histogram of oy. its central value of «13 km s! is a goo estimate of the standard error of a single velocity measurement based on that line.," By analogy with the histogram of $\sigma_{\rm H}$, its central value of $\sim$ 13 km $^{-1}$ is a good estimate of the standard error of a single velocity measurement based on that line."
+ It is slightly smaller than Avy because hydrogen lines are very broad in most objects. which results 11 increased fitting errors.," It is slightly smaller than $\Delta v_{\rm H}$ because hydrogen lines are very broad in most objects, which results in increased fitting errors."
+ Also. the wavelength calibration ts o1 the average sligthly less accurate at Hj than at the Na doublet because there are fewer lamp lines in the bluer part of the spectrum.," Also, the wavelength calibration is on the average sligthly less accurate at $_\beta$ than at the Na doublet because there are fewer lamp lines in the bluer part of the spectrum."
+ Velocity measurement from magnesium lines was possible for only 39 objects., Velocity measurement from magnesium lines was possible for only 39 objects.
+ In nearly all of them the best visible was Mel 5183.6A., In nearly all of them the best visible was MgI 5183.6.
+. Unfortunately. for some objects even that line could not be well fitted in all spectra. so that their vy and ome had to be calculated from as few as four spectra.," Unfortunately, for some objects even that line could not be well fitted in all spectra, so that their $\bar v_{\rm Mg}$ and $\sigma_{\rm Mg}$ had to be calculated from as few as four spectra."
+ As a result. the number of chi-square degrees of freedom was not well determined for the oy. observable.," As a result, the number of chi-square degrees of freedom was not well determined for the $\sigma_{\rm Mg}$ observable."
+ Because of that. the histogram of cw in Fig.," Because of that, the histogram of $\sigma_{\rm Mg}$ in Fig."
+ 5 is not directly comparable to the remaining two histograms., \ref{fig:sigHb_sigNa} is not directly comparable to the remaining two histograms.
+ It was plotted for illustrative purpose only. and the rms deviations oy. were not used for further analysis.," It was plotted for illustrative purpose only, and the rms deviations $\sigma_{\rm Mg}$ were not used for further analysis."
+ We only used the mean velocities Ty. to check if they differ from their fy counterparts., We only used the mean velocities $\bar v_{\rm Mg}$ to check if they differ from their $\bar v_{\rm H}$ counterparts.
+ No systematic differences were found. and the mean difference was equal to just 5.3 km s!. which confirmed the reliability of velocity measurements based on Hj.," No systematic differences were found, and the mean difference was equal to just 5.3 km $^{-1}$, which confirmed the reliability of mean-velocity measurements based on ${_\beta}$."
+" The FXCOR measurements fulfilled the reliability criterion [7>0.2 in the case of 28 objects. while the stellar Na D, was strong enough for successful fitting in only 11 objects."," The FXCOR measurements fulfilled the reliability criterion $f_c^m>0.2$ in the case of 28 objects, while the stellar Na $_1$ was strong enough for successful fitting in only 11 objects."
+ As a result. there were Just 9 objects for which we collected four complete independent velocity measurements (from Hj. Mg and stellar Na fitting. and from FXCOR).," As a result, there were just 9 objects for which we collected four complete independent velocity measurements (from $_\beta$, Mg and stellar Na fitting, and from FXCOR)."
+ For each spectrum of those objects we calculated the average and the corresponding rms deviation oy., For each spectrum of those objects we calculated the average and the corresponding rms deviation $\sigma_4$.
+ Fig., Fig.
+" 6 shows that there are no significant differences between v, and velocities obtained from Πρ fitting. thus proving the reliability of the latter."," \ref{fig:errHb_err3} shows that there are no significant differences between $v_4$ and velocities obtained from $_\beta$ fitting, thus proving the reliability of the latter."
+ All deviationsσι except one range between 2 and 20 kms., All deviations$\sigma_4$ except one range between 2 and 20 km $^{-1}$ .
+ We note that FXCOR returnsformal velocity errors. but they are correct only to within a scaling factor which depends on the number of counts in the spectra and the Fourier filter parameters used (seee.q.Foretal. 2010)..," We note that FXCOR returnsformal velocity errors, but they are correct only to within a scaling factor which depends on the number of counts in the spectra and the Fourier filter parameters used \citep[see e.q.][]{for10}. ."
+ For most of our, For most of our
+Alassive WR-OB star binaries may have 210 keV X-ray luminosities up toI0ere|. presumably. due to wind-wind collisions (Portegies. Zwart. Pooley Lewin 2002).,"Massive WR-OB star binaries may have 2–10 keV X-ray luminosities up to$10^{35}$, presumably due to wind-wind collisions (Portegies Zwart, Pooley Lewin 2002)."
+ 1n addition they often show a strong iron I-line with an equivalent width of ~1 keV. However. here the. problem. is the factor LOO change in luminosity for exiunple. an extensive study covering the dillerent orbital phases shows he variability in the hard X-ray component in such systenis is ato most a [actor of three (Alaeda 1900).," In addition they often show a strong iron K-line with an equivalent width of $\sim$ 1 keV. However, here the problem is the factor 100 change in luminosity – for example, an extensive study covering the different orbital phases shows the variability in the hard X-ray component in such systems is at most a factor of three (Maeda 1999)."
+ A rackerouncd active galaxy. seen hrough the Galaxy. might conceivably be a candidate but again it is very cdillicult o match the observations with any known class of object.," A background active galaxy, seen through the Galaxy, might conceivably be a candidate but again it is very difficult to match the observations with any known class of object."
+" For example. in Sevlert Lo nuclei the measured 6.7-keV line equivalent width is invariably less than a few hundred eV. Nandra Pounds. 1994: Reeves ""Turner. 2000). whereas the extreme variability exeludes an obscurecl Seyfert LL nucleus IExovama 1989) as a potential counterpart."," For example, in Seyfert I nuclei the measured 6.7-keV line equivalent width is invariably less than a few hundred eV, Nandra Pounds 1994; Reeves Turner 2000), whereas the extreme variability excludes an obscured Seyfert II nucleus Koyama 1989) as a potential counterpart."
+ The nature of2913.0 is remarkably similar to AN 0423. which was discovered. in an survey of the Scutum arm region in 1996 (Lerada 1999).," The nature of is remarkably similar to AX $-$ 0423, which was discovered in an survey of the Scutum arm region in 1996 (Terada 1999)."
+ Terada (1999) found that AX 0423 has a strong iron line at 6.8 keV with an equivalent. width of ~4 keV. Phe spectrum is well approximated with a thin-thermal plasma model witha temperature of Pt keV ancl heavy-metal abundance of 73 solar., Terada (1999) found that AX $-$ 0423 has a strong iron line at 6.8 keV with an equivalent width of $\sim$ 4 keV. The spectrum is well approximated with a thin-thermal plasma model witha temperature of $\sim$ 4 keV and heavy-metal abundance of $\sim$ 3 solar.
+ AX 0423 also showed transient behaviour on a time-scale of less than hall avear., AX $-$ 0423 also showed transient behaviour on a time-scale of less than half a year.
+" In adedition there are two other sources which have an extremely strong iron line: AX 592 (CP ""Tuc: Misaki 1996) and RN 1802.1|1804 (VSSA Ler: Ishida 1998).", In addition there are two other sources which have an extremely strong iron line: AX $-$ 592 (CP Tuc; Misaki 1996) and RX J1802.1+1804 (V884 Her; Ishida 1998).
+ Both of the latter sources have been identified: as polars AM Her typeCVs)., Both of the latter sources have been identified as polars AM Her typeCVs).
+ Terada (1999. 2001) concluded. that the three," Terada (1999, 2001) concluded that the three"
+be.,be.
+ Π the lar 3-kpe ari sample of sources was significantly biased by distant sources. we would expect a shallow distribution of peak flux densities. whilst if it was biased by nearby sources. we would expect au asyiumetric cistribution with a peak at very low peak flux deusities.," If the far 3–kpc arm sample of sources was significantly biased by distant sources, we would expect a shallow distribution of peak flux densities, whilst if it was biased by nearby sources, we would expect an asymmetric distribution with a peak at very low peak flux densities."
+ We do not see either. instead findiug ai approxtinately Ciaussian distribution with a median of 2.1 Jy. unplviug a population dominated by intermediate distances.," We do not see either, instead finding an approximately Gaussian distribution with a median of 2.4 Jy, implying a population dominated by intermediate distances."
+ Furthermore. the [ar spiral arms are not well traced by or CO at these longitudes. so their location is less precise.," Furthermore, the far spiral arms are not well traced by or CO at these longitudes, so their location is less precise."
+ Regious of higl-uiass star formation are comparatively rare. so without either of the spiral armis being tangential. we would uot expect there to be a significant population present in either.," Regions of high-mass star formation are comparatively rare, so without either of the spiral arms being tangential, we would not expect there to be a significant population present in either."
+ These factors combine to give us coulicdence that the dense region iceutified by ? aud ? is a product of lar 3-kpe arm sources interacting with the bar., These factors combine to give us confidence that the dense region identified by \citet{caswell10mmb1} and \citet{green10mmb2} is a product of far 3–kpc arm sources interacting with the bar.
+ Expaucdiug on the analysis of 2.. we assume a bar orientation of 15° (e.g.?).. with the axis left as a [ree parameter.," Expanding on the analysis of \citet{green10mmb2}, we assume a bar orientation of $^{\circ}$ \citep[e.g.][]{benjamin05}, with the semi-major axis left as a free parameter."
+ The deusity enhancement at the far end of the bar (between lougitudes —97 aud —157) thus implies a semi-major axis of the bar between 2.2 aud. 13 kpc., The density enhancement at the far end of the bar (between longitudes $-$ $^{\circ}$ and $-$ $^{\circ}$ ) thus implies a semi-major axis of the bar between 2.2 and 4.3 kpc.
+ This iudicates that the near end of the bar lies between loungitudes 137 and 307., This indicates that the near end of the bar lies between longitudes $^{\circ}$ and $^{\circ}$.
+" We see from Table 1 aud Figure 3. that there are three higher density regions of ruasers within this range: betweenLO’ aud 11 (1 bin above 6055::5,): between IS"" and 22° longitude (1 bin above 2655:::,55): aud between 21° and 28° longitude (1 bin above 26,5,5::45,): these correspond to semianajor axes of 1.8—2.kkpe. kkpe and 3.7-Likkpe respectively."," We see from Table \ref{resotable} and Figure \ref{lvdensity} that there are three higher density regions of masers within this range: between$^{\circ}$ and $^{\circ}$ (1 bin above $\sigma_{poisson}$ ); between $^{\circ}$ and $^{\circ}$ longitude (1 bin above $\sigma_{poisson}$ ); and between $^{\circ}$ and $^{\circ}$ longitude (1 bin above $\sigma_{poisson}$ ); these correspond to semi-major axes of kpc, kpc and kpc respectively."
+ In dyuamical simulations. bars with short semi-major axes (or the asstunptiou of such a mass distribution) are almost exclusively associated with small ji Orientation angles (e.g.2???) aud a short bar with au acute orientation augle would produce densities in the far 3-kpce arm at longitudes sinaller than is observed (?)..," In dynamical simulations, bars with short semi-major axes (or the assumption of such a mass distribution) are almost exclusively associated with small bar orientation angles \citep[e.g.][]{binney91,freudenreich98, babusiaux05} and a short bar with an acute orientation angle would produce densities in the far 3–kpc arm at longitudes smaller than is observed \citep{green10mmb2}."
+ Additionally. a shorter oir of this type is associated with the Galactic bulge. but the 6.77CHz inethiauol maser population IHs a Barrow latitude distribution unassociated with the bulge (?)..," Additionally, a shorter bar of this type is associated with the Galactic bulge, but the 6.7–GHz methanol maser population has a narrow latitude distribution unassociated with the bulge \citep{caswell10mmb1}."
+ Heuce the semi-major axis is uilikely to be in the range 1.5—2.E kpc., Hence the semi-major axis is unlikely to be in the range 1.8–2.4 kpc.
+ The estimate of a raclius of corotation resonance at  Lkkpe implies that a semi-major axis of kkpe is also unlikely., The estimate of a radius of corotation resonance at $\sim$ kpc implies that a semi-major axis of kpc is also unlikely.
+ IL. ou the other παπα. we assume the semi-major axis of the bar is fixed at kkpc (the estimate of ?/— and references therein. scaled to 8.1 kpc) aud the bar orientation is iustead left as a [ree parameter. the deusity enliaucement at the far eid of the bar implies a bar orientation between 35° aud 53°.," If, on the other hand, we assume the semi-major axis of the bar is fixed at kpc (the estimate of \citealt{gerhard02} and references therein, scaled to 8.4 kpc) and the bar orientation is instead left as a free parameter, the density enhancement at the far end of the bar implies a bar orientation between $^{\circ}$ and $^{\circ}$ ."
+ This rauge of orientation augles would locate the near end of the bar between lougitudes 20° aud 21°. which overlaps with the high density of masers seen at velocities close to ! (1 bin above poisson 1n Figure 3)).," This range of orientation angles would locate the near end of the bar between longitudes $^{\circ}$ and $^{\circ}$, which overlaps with the high density of masers seen at velocities close to $^{-1}$ (1 bin above $\sigma_{poisson}$ in Figure \ref{lvdensity}) )."
+ Iu summary we believe the maser population traces the influence of a long thin bar with a axis of 3.I kpe aud au orientatiou of  15°., In summary we believe the maser population traces the influence of a long thin bar with a semi-major axis of $\sim$ 3.4 kpc and an orientation of $\sim$ $^{\circ}$ .
+ This implies that. if both a lougaud short bar exist within our Galaxy. it is the long component which is primarily traceable by (higli-1uass) star," This implies that, if both a longand short bar exist within our Galaxy, it is the long component which is primarily traceable by (high-mass) star"
+One of the most important. aim of modern extragalactic astronomy ancl cosmology is lo solve (he problem of structure formation.,One of the most important aim of modern extragalactic astronomy and cosmology is to solve the problem of structure formation.
+ There are many theories used (o develop scenarios of structure formations (Peebles1969:Zeldovich1970:Sunvaew&1972:Doroshkevieh1973:Shandarin1974:Dekel19385:Wesson1982:Silk1983:Bower 2005).," There are many theories used to develop scenarios of structure formations \citep{Peebles69,Zeldovich70,Sunyaew72,Doroshkevich73,Shandarin74,Dekel85,Wesson82,Silk83,Bower05}."
+. In the commonly accepted. ACDM moclel. the Universe deems (to be spatially flat. as well as," In the commonly accepted $\Lambda$ CDM model, the Universe deems to be spatially flat, as well as"
+"statistic to measure the substructure or “lumpiness” that we designate oym(r), which is created by an unsharp- procedure: where RH, is the effective (half-mass) radius of the galaxy, p is the stellar density field (created by cloud-in-cell mapping of the star particles to a grid of resolution 0.38 kpc, or twice the gas softening length), and G(r) is a gaussian of width r.","statistic to measure the substructure or “lumpiness” that we designate $\sigma_{\text{UM}}(r)$, which is created by an unsharp-mask-like procedure: where $R_e$ is the effective (half-mass) radius of the galaxy, $\rho$ is the stellar density field (created by cloud-in-cell mapping of the star particles to a grid of resolution 0.38 kpc, or twice the gas softening length), and $G(r)$ is a gaussian of width $r$."
+" That is, we create a mask by smoothing the stellar density field of the galaxy with a fixed-width gaussian kernel, then sum ((data—mask)/mask)? over all pixels in the grid where it is positive (i.e. overdense regions) except the central peak, and normalize to the size of the galaxy."," That is, we create a mask by smoothing the stellar density field of the galaxy with a fixed-width gaussian kernel, then sum $((\text{data}-\text{mask})/\text{mask})^2$ over all pixels in the grid where it is positive (i.e. overdense regions) except the central peak, and normalize to the size of the galaxy."
+" We find that the Old UV model indeed has far more substructure than the two new models, at least for halos A and C; Table 4 shows cuw(4kpc) for the three backgrounds and ICs."," We find that the Old UV model indeed has far more substructure than the two new models, at least for halos A and C; Table \ref{tab:substruct} shows $\sigma_{\text{UM}}(4\text{kpc})$ for the three backgrounds and ICs."
+" Halo E, as remarked on in is composed mainly of stars which formed in- (2007),,rather than the accretion of smaller stellar systems, and therefore we expect its substructure to be much less affected by the ionizing background, which is what we find."," Halo E, as remarked on in, is composed mainly of stars which formed in-situ rather than the accretion of smaller stellar systems, and therefore we expect its substructure to be much less affected by the ionizing background, which is what we find."
+" At first glance, Fig."," At first glance, Fig."
+" 8 and Table 4 may seem to be incompatible, since the two models without X-rays are nearly the same in the former and vastly different in the latter."," \ref{fig:haloint-A100} and Table \ref{tab:substruct} may seem to be incompatible, since the two models without X-rays are nearly the same in the former and vastly different in the latter."
+" However, if we increase the outer radius in the definition of oy» from two to three effective radii (roughly from 20 to 30 kpc for Galaxy A), we find that becomes (125.7,145.8, for Old UV, New UV, and oum(4kpc)New UV-+X respectively, a 31.5)result much more in accordance with Fig. 8.."," However, if we increase the outer radius in the definition of $\sigma_{\text{UM}}$ from two to three effective radii (roughly from 20 to 30 kpc for Galaxy A), we find that $\sigma_{\text{UM}}(4\text{kpc})$ becomes $(125.7, 145.8, 31.5)$ for Old UV, New UV, and New UV+X respectively, a result much more in accordance with Fig. \ref{fig:haloint-A100}."
+" That is, the New UV simulation has numerous small halos, but they are at larger radii from the central galaxy than in the Old UV case."," That is, the New UV simulation has numerous small halos, but they are at larger radii from the central galaxy than in the Old UV case."
+" On the other hand, Old UV and New UV have the same number of independent halos at the present, as seen Fig. 8,,"," On the other hand, Old UV and New UV have the same number of independent halos at the present, as seen Fig. \ref{fig:haloint-A100},"
+ but Old UV has significantly more substructure., but Old UV has significantly more substructure.
+" This means that Old UV formed more small halos initially, but they were accreted onto the central galaxy by z~0.5, enhancing its substructure."," This means that Old UV formed more small halos initially, but they were accreted onto the central galaxy by $z\simeq0.5$, enhancing its substructure."
+ Figure 10 shows contours of stellar density for the central 20kpc of galaxy A with the four background models., Figure \ref{fig:starcont} shows contours of stellar density for the central 20kpc of galaxy A with the four background models.
+" The smaller size of New UV+X is readily apparent, as is the reduction in the number and mass of subhalos."," The smaller size of New UV+X is readily apparent, as is the reduction in the number and mass of subhalos."
+ One can also see the sharp central peak in New UV and New UV+X which is absent in Old UV., One can also see the sharp central peak in New UV and New UV+X which is absent in Old UV.
+" 'Taken collectively, our results suggest a picture of the effects of an increased early UV radiation background along the following lines."," Taken collectively, our results suggest a picture of the effects of an increased early UV radiation background along the following lines."
+" At early times (z=2), the more intense/harder radiation and earlier reionization make the gas hotter, especially the less-dense gas, which cannot cool as effectively (Fig. 2))."," At early times $z\gtrsim 2$ ), the more intense/harder radiation and earlier reionization make the gas hotter, especially the less-dense gas, which cannot cool as effectively (Fig. \ref{fig:Thist-all}) )."
+" That is to say, the processes of H and He reionization create large injections of heat, so even though the gas begins to cool one reionization is complete, an earlier reionization means the gas spends less time in the cold, neutral state it has been in since recombination."," That is to say, the processes of H and He reionization create large injections of heat, so even though the gas begins to cool one reionization is complete, an earlier reionization means the gas spends less time in the cold, neutral state it has been in since recombination."
+" At very early times (z>4) the gas has not yet equilibrated, and so the more intense but softer FG UV background gives a higher temperature than Old UV, but by z—3.2 the situation has reversed, consistent with."," At very early times $z>4$ ) the gas has not yet equilibrated, and so the more intense but softer FG UV background gives a higher temperature than Old UV, but by $z=3.2$ the situation has reversed, consistent with."
+" At late times, this means that the primary galaxy has less substructure, since there are fewer small stellar systems to accrete; it contains more gas in its central regions (Fig. 4)),"," At late times, this means that the primary galaxy has less substructure, since there are fewer small stellar systems to accrete; it contains more gas in its central regions (Fig. \ref{fig:cgasacc-A100}) ),"
+" and hence more late in-situ star formation (Fig. 6)),"," and hence more late in-situ star formation (Fig. \ref{fig:sfr-gal-A100}) ),"
+" since that gas would otherwise have formed stars before falling in; it is still smaller but now more tightly bound due to gaining less energy from dynamical friction and more central star (Naab,formation from infalling cold streams.", since that gas would otherwise have formed stars before falling in; it is still smaller but now more tightly bound due to gaining less energy from dynamical friction and more central star formation from infalling cold streams.
+" In the terms of the picture, the cold gas inflows persist longer when early (2007)star formation is suppressed by ionizing radiation."," In the terms of the picture, the cold gas inflows persist longer when early star formation is suppressed by ionizing radiation."
+" We can distinguish the effects of intensity and spectral shape by comparing New UV and FG UV, since they have roughly the same intensity (H photoionization rate, etc.)"," We can distinguish the effects of intensity and spectral shape by comparing New UV and FG UV, since they have roughly the same intensity (H photoionization rate, etc.)"
+ from 2«z8 but FG UV has a significantly softer spectrum due to the rapid falloff of the quasar contribution., from $2<z<8$ but FG UV has a significantly softer spectrum due to the rapid falloff of the quasar contribution.
+" We find that FG UV in general gives results intermediate between Old UV and New UV, for gas temperature, rate of gas accretion by the central galaxy, star-formation rate, and concentration of stellar mass; in other words, an increase in the intensity of early radiation and an increase in the hardness of the background spectrum produce qualitatively the same results in most"," We find that FG UV in general gives results intermediate between Old UV and New UV, for gas temperature, rate of gas accretion by the central galaxy, star-formation rate, and concentration of stellar mass; in other words, an increase in the intensity of early radiation and an increase in the hardness of the background spectrum produce qualitatively the same results in most"
+"line, the large","line, the large"
+Note that au instantaneous spectra that is flat Voxe (c.g.Ixatzetal.2007) or harder (eg.Ieshet2006:Steckeretal.2007.iuxothers) can Ὡς> obtained -iu specific particle acceleration models.,"Note that an instantaneous spectrum that is flat $N\propto \vep^{-2}$ \citep[e.g.][]{Katz07} or harder \citep[e.g.][and others]{Keshet06, Stecker07} can be obtained in specific particle acceleration models."
+ These models. however. are ditffereut frou the possible hard spectrin of escaping particles discussed above.," These models, however, are different from the possible hard spectrum of escaping particles discussed above."
+ In fact. such hare Instantaneous spectra lead to a flat escaped spectrin Nate)x©2 Jsee discussion followine eq. (5)]].," In fact, such hard instantaneous spectra lead to a flat escaped spectrum $\Nesc(\vep)\propto \vep^{-2}$ [see discussion following eq. ]."
+ The iuportant point is that the escaping spectraids highly uncertain. very sensitive to the acceleration mechamisin. aud may be cousicerably harder than ο)XTE27.," The important point is that the escaping spectrum is highly uncertain, very sensitive to the acceleration mechanism, and may be considerably harder than $\Nesc(\vep)\propto \vep^{-2}$."
+ Ounce he shock decelerates to nonu relativistic speeds. the iuteeratede escaping spectrmm changes its form.," Once the shock decelerates to non relativistic speeds, the integrated escaping spectrum changes its form."
+ Iu the subsequent Sedov-Tavlor phase. 1ο ininimal CR energy does not change any more. £io6? aud the escaping spectrum will be similar to the insautaneous spectrim [see Eq. (6)]].," In the subsequent Sedov-Taylor phase, the minimal CR energy does not change any more, $\vep_{\min}\sim m_pc^2$, and the escaping spectrum will be similar to the instantaneous spectrum [see Eq. ]."
+" Note that iu the exτοις Case in which CRs carry most of the euerev. aux assunine they are accelerated bv nou-linear DSA (forarecentreviewseeMalkov&Drury 2001).. the iusautaneous spectruni nav not be a power law and correspondingly the escapiug spectrum is differeut than that o€""ven here (οιο,Capriolietal.2009:Ohira2009)."," Note that in the extreme case in which CRs carry most of the energy, and assuming they are accelerated by non-linear DSA \citep[for a recent review see][]{Malkov01}, the instantaneous spectrum may not be a power law and correspondingly the escaping spectrum is different than that given here \citep[e.g.][]{Caprioli09,Ohira09}."
+". Tn anv case. a xenificaut break is expected in the escaping spectrin at an cherey z, equal to the maximal euergv acüievable at the transition from relativistic to non-rolativistic shock velocities."," In any case, a significant break is expected in the escaping spectrum at an energy $\vep_\ast$ equal to the maximal energy achievable at the transition from relativistic to non-relativistic shock velocities."
+ Using Eq., Using Eq.
+ for Po—1. the break is expected at where we have taken. as al example. a gauunneraxy burst source with E—IUEzerg. p=ngny is the züubieut deusitv with v=nycem?. and eg=O.lep4.," for $\Gamma\bt\sim1$, the break is expected at where we have taken, as an example, a gamma-ray burst source with $E=10^{51}E_{51}\erg$, $\rho=n~m_p$ is the ambient density with $n=n_0\cm^{-3}$, and $\epsilon_B=0.1\epsilon_{B,-1}$."
+ The spectral index above this energy willbe 2 . as given by Eq.," The spectral index above this energy will be $-2-\xesc=-2+5x$ , as given by Eq."
+(10)... We note that the energy range where suus=fe is limited between s. aud Πακ)~zyax(Lo) where Dy is the initial Loreutz factor., We note that the energy range where $\xesc=-5x$ is limited between $\vep_\ast$ and $\max(\vep)\sim\vep_{\max}(\Gamma_0)$ where $\Gamma_0$ is the initial Lorentz factor.
+ The maximal cucrey scales with D as ΌμωςxIBD resulting in a narrow energy range 2.S9XDO, The maximal energy scales with $\Gamma$ as $\vep_{\max}\propto\Gamma^{1/3}$ resulting in a narrow energy range $\vep_\ast\lesssim\vep\lesssim\Gamma_0^{1/3}\vep_\ast$.
+ While this range depends oi model parameters. it docs point to a possible πουίσια for achieving a hardening of the spectrum at the lieh eud of the CR spectiua.," While this range depends on model parameters, it does point to a possible mechanism for achieving a hardening of the spectrum at the high end of the CR spectrum."
+ We conclude that the spectrmm of UITECRs. may be different than 27. aud possibly considerably harder. e.g. if they originate from a relativistic decelerating blast- resulting from the explosion of a stellar mass object.," We conclude that the spectrum of UHECRs, may be different than $\vep^{-2}$, and possibly considerably harder, e.g. if they originate from a relativistic decelerating blast-wave resulting from the explosion of a stellar mass object."
+ Dk EW acknowledge support by ISF. AEC aud Minerva grants;," BK EW acknowledge support by ISF, AEC and Minerva grants."
+ PAL acknowledges support from NSF PIIY-0757155 erant., PM acknowledges support from NSF PHY-0757155 grant.
+We eive here an explicit expression for the d-mode in an homogeneous but semi-infinite medium.,We give here an explicit expression for the d-mode in an homogeneous but semi-infinite medium.
+" The d-mocle is the solution of the diffusion equation where (he speed w is assumed constant because of homogeneity,", The d-mode is the solution of the diffusion equation where the speed $u$ is assumed constant because of homogeneity.
+ We solve first the equation wilh u=0: to do so. we remark (hat suitable boundary conditions ave that df—0 asc—cox. depending on whether we are considering downstream or upstream regions. respectively.," We solve first the equation with $u = 0$; to do so, we remark that suitable boundary conditions are that $\delta\! f\rightarrow 0$ as $x\rightarrow \pm\infty$, depending on whether we are considering downstream or upstream regions, respectively."
+" The solution can be obtained by separation of variables. obtaining: subject to the constraint The sign of &, is the one that allows the solution to remain finite at infinity."," The solution can be obtained by separation of variables, obtaining: subject to the constraint The sign of $k_x$ is the one that allows the solution to remain finite at infinity."
+" Solutions belonging to different. values of A, and p can obviously be superposed. but. we know that"," Solutions belonging to different values of $k_x$ and $p$ can obviously be superposed, but we know that"
+mav be responsible for the L to T transition.,may be responsible for the L to T transition.
+ Durgasserοἱal.(2002) tested this hypothesis with a simple ‘tov model by summiug weighted contributions of the spectra of cloudy and cloudless models., \citet{Bur02} tested this hypothesis with a simple `toy model' by summing weighted contributions of the spectra of cloudy and cloudless models.
+ Thev showed that the observed J band brightening across the transition could arise lrom decreasing cloud coverage., They showed that the observed $J$ band brightening across the transition could arise from decreasing cloud coverage.
+ Further support for this hypothesis comes from observations of L and T dwarf variability., Further support for this hypothesis comes from observations of L and T dwarf variability.
+ While previous studies were somewhat equivocal (simmarized in Artigauetal. (2009))). two early T dwarls have recently been shown to have large near-infrared photometric variability (Artigauetal.2009:BRadigan2010) consistent with surface variations in cloud coverage modulated by rotation.," While previous studies were somewhat equivocal (summarized in \citet{Art09}) ), two early T dwarfs have recently been shown to have large near-infrared photometric variability \citep{Art09,Rad10} consistent with surface variations in cloud coverage modulated by rotation."
+ The approach to modeling holes of Durgasser et wwas hiehlv simplistic., The approach to modeling holes of Burgasser et was highly simplistic.
+ The principal shortcoming being that it is not physically correct to combine the contributions of separate cloudy. ancl cloudless models., The principal shortcoming being that it is not physically correct to combine the contributions of separate cloudy and cloudless models.
+ Deep in the atmosphere of a brown dwarf the entropy in the convection zone must match that of the interior., Deep in the atmosphere of a brown dwarf the entropy in the convection zone must match that of the interior.
+ Thus the temperature al a given. deep. pressure level is expected (o be horizontally constant.," Thus the temperature at a given, deep, pressure level is expected to be horizontally constant."
+ However for a fixed Ziy a cloudxy atmosphere is evervwhere hotter than a cloudless atmosphere., However for a fixed $\teff$ a cloudy atmosphere is everywhere hotter than a cloudless atmosphere.
+ As an example Figure 1 presents model atmosphere profiles lor a unilorm cloudy and a cloudless atmosphere following the techniques of Marleyοἱal.(2002) ancl Saumon&Marley(2008)., As an example Figure 1 presents model atmosphere profiles for a uniform cloudy and a cloudless atmosphere following the techniques of \citet{Mar02} and \citet{Sau08}.
+. At depth. the cloudy profile is warmer than the cloudless profile bv over Ix: the dillerence in some models is even greater.," At depth, the cloudy profile is warmer than the cloudless profile by over $\,$ K; the difference in some models is even greater."
+ Thus standard. cloudy and clouclless models cannot simultaneously be valid descriptions of the real atmosphere at (wo locations even though both models are descriptions of an atmosphere with the same Z;r., Thus standard cloudy and cloudless models cannot simultaneously be valid descriptions of the real atmosphere at two locations even though both models are descriptions of an atmosphere with the same $\teff$.
+ Clearly a new technique for sell-consistentlv. treating partly cloudy atmospheres is reeuired., Clearly a new technique for self-consistently treating partly cloudy atmospheres is required.
+ While we do not vet understand. why cloud holes might appear. we here present a new approach inspired from models of Earth's atmosphere to model their influence and apply our results to model the spectra ancl colors of L ancl T cdwarfs.," While we do not yet understand why cloud holes might appear, we here present a new approach inspired from models of Earth's atmosphere to model their influence and apply our results to model the spectra and colors of L and T dwarfs."
+ ]Iustead of combiningseperately computed. profiles for purely cloudy and cloud-free clwarls we wish to construct aglobal temperature-pressure profile £C?) Chat incorporates simultaneously the influences of both cloudy and eloud-Iree regions on the energv balance of the atmosphere., Instead of combining computed profiles for purely cloudy and cloud-free dwarfs we wish to construct a temperature-pressure profile $T(P)$ that incorporates simultaneously the influences of both cloudy and cloud-free regions on the energy balance of the atmosphere.
+ The final profile should conserve the total [αν while allowing for nearby ablmospheric regions to have differing cloul — but not thermal — profiles., The final profile should conserve the total flux while allowing for nearby atmospheric regions to have differing cloud – but not thermal – profiles.
+ This conceptually allows clouds to be displaced by winds. updralts. or downdralts and change location. as long as (he global mean cloud fraction is constant.," This conceptually allows clouds to be displaced by winds, updrafts, or downdrafts and change location, as long as the global mean cloud fraction is constant."
+ In (hree-dimensional terrestrial numerical weather prediction and climate models. clouds are (vpically smaller than the adopted computational erid scale.," In three-dimensional terrestrial numerical weather prediction and climate models, clouds are typically smaller than the adopted computational grid scale."
+ Various methods are used, Various methods are used
+the first episode of mass transfer must proceed in a dynamically stable iinuer.,the first episode of mass transfer must proceed in a dynamically stable manner.
+ Whether. ouce started. ALT will proceed in a stable or unstable manner depends ou the response of the donor and its Roche lobe to the AIT.," Whether, once started, MT will proceed in a stable or unstable manner depends on the response of the donor and its Roche lobe to the MT."
+ Mass loss is a perturbation that brings a star out of lvdrostatic and thermal equilibrium., Mass loss is a perturbation that brings a star out of hydrostatic and thermal equilibrium.
+ In order to re-establish these equilibria. the star will either expaud or coutract. first restoring hwdrostatie equilibriuu. aud then. ou a lounger timescale. thermal equilibrimu.," In order to re-establish these equilibria, the star will either expand or contract, first restoring hydrostatic equilibrium, and then, on a longer timescale, thermal equilibrium."
+ Like the stars radius. its Roche-lobe radius also changes iu response to the MT.," Like the star's radius, its Roche-lobe radius also changes in response to the MT."
+ As a stabilitv condition. we thus demand that upon AIT the new donor radius should remain within the donors new Roche lobe.," As a stability condition, we thus demand that upon MT the new donor radius should remain within the donor's new Roche lobe."
+ The stability of the AIT depends on the response of the stellar radius to mass loss., The stability of the MT depends on the response of the stellar radius to mass loss.
+ To consider stability in detail. we will use the Luear stability analysis following ??..," To consider stability in detail, we will use the linear stability analysis following \citet{HW87, SPV97}."
+ We define the adiabatic massradius exponent (the donors response to the mass loss on au adiabatic timescale) as and the Roche lobes massradius exponent (the Roche-lobe respouse to the MT) as The criterion for dvuamical stability of AIT then cau be written as Qaq=GRL-, We define the adiabatic mass–radius exponent (the donor's response to the mass loss on an adiabatic timescale) as and the Roche lobe's mass–radius exponent (the Roche-lobe response to the MT) as The criterion for dynamical stability of MT then can be written as $\zeta_\mathrm{ad} \ge \zeta_\mathrm{RL}$.
+ If this criterion is satisfied. the donor is able to recover its hydrodynamical equilibrium while remaiuiug within its Roche lobe.," If this criterion is satisfied, the donor is able to recover its hydrodynamical equilibrium while remaining within its Roche lobe."
+ It will then try to recover its thermal equilibrium on the Welwin-Uehuholtz timescale., It will then try to recover its thermal equilibrium on the Kelvin-Helmholtz timescale.
+ The chauge in its equilibriu radius cau be expressed with the use of the equilibrimm massradius expoucut Tf. in additiou to the dvuamicalstability couditiou. the condition 2Qgp also holds. thea MIT will proceed onu the uuclear-evolutionQa; timescale i a secularly stable manner.," The change in its equilibrium radius can be expressed with the use of the equilibrium mass–radius exponent If, in addition to the dynamical-stability condition, the condition $\zeta_\mathrm{eq} \ge \zeta_\mathrm{RL}$ also holds, then MT will proceed on the nuclear-evolution timescale in a secularly stable manner."
+ Tf QucoORLcGeqe MT will be driven by thermal readjustiuent auc will proceed stably ou the thermal timescale (524).," If $\zeta_\mathrm{ad} \ge \zeta_\mathrm{RL} > \zeta_\mathrm{eq}$, MT will be driven by thermal readjustment and will proceed stably on the thermal timescale $\tau_\mathrm{th}$ )."
+ In the case of a red eiaut. the radius of the star iu complete equilibriu is a function predominantly of its core mass. aud as such its thermal response is usually adopted to be coq20. though it can vary.," In the case of a red giant, the radius of the star in complete equilibrium is a function predominantly of its core mass, and as such its thermal response is usually adopted to be $\zeta_\mathrm{eq} \approx 0$, though it can vary."
+ We will now consider these respouses iu more detail., We will now consider these responses in more detail.
+ The adiabatic massradius exponent for giants can be obtained considering the case of a condeused polvtrope with »=3/2 (?).., The adiabatic mass–radius exponent for giants can be obtained considering the case of a condensed polytrope with $n=3/2$ \citep{HW87}.
+" The value of σα cau then be fond to wi accuracy usine the formula from ?:: where AL,=om,(AL.)fin is the mass fraction of the helimm core.", The value of $\zeta_\mathrm{ad}$ can then be found to within accuracy using the formula from \citet{SPV97}: where $M_\mathrm{c} \equiv m_\mathrm{c}(M_\odot)/m$ is the mass fraction of the helium core.
+ We expand the logithnuüe derivative of the orbital separation with respect to the donor mass as: ere q=maf£ma is the mass ratio., We expand the logarithmic derivative of the orbital separation with respect to the donor mass as: Here $q=m_\mathrm{d}/m_\mathrm{a}$ is the mass ratio.
+" The first ται, OluefOhimn.1is solely due to the mass loss or ALT aud can be found using eq.(L1)): We will be most concerned with stability at the ouset of a fast initial AIT phase (see 8L3. 55.3)."," The first term, $\partial \ln a/ \partial \ln m_\mathrm{}$, is solely due to the mass loss or MT and can be found using \ref{orbit}) ): We will be most concerned with stability at the onset of a fast initial MT phase (see $\S 4.3$, $\S5.3$ )."
+ In this case we lav assunie that the AIT rate from the donor exceeds the wind mass-loss rate significantly. so that for the purpose of this stability analvsis we cau assiuue that 6=0 (See 2.2)).," In this case we may assume that the MT rate from the donor exceeds the wind mass-loss rate significantly, so that for the purpose of this stability analysis we can assume that $\delta=0$ (See \ref{sec:2.2}) )."
+ Then: where J ds the fraction of the transferred. mass that is accreted by the secoudary star., Then: where $\beta$ is the fraction of the transferred mass that is accreted by the secondary star.
+ We note that the presence of wind makes MT iore stable. as the orbit always expands in response to wind loss.," We note that the presence of wind makes MT more stable, as the orbit always expands in response to wind loss."
+ As such. with à=0 we would obtain a stricter criterion for MUT stability.," As such, with $\delta=0$ we would obtain a stricter criterion for MT stability."
+ The above equation demonstrates that the Roche-lobe response is a function of the MT couservation factor: iu eeneral Ganτομνα).," The above equation demonstrates that the Roche-lobe response is a function of the MT conservation factor; in general, $\zeta_\mathrm{RL} \approx \zeta_\mathrm{RL}(\beta \mathrm{,q})$."
+" The second term in refeqiuzeta,JHeonsistsoftheBtochelobe sresponsctotheehaugeiimassrc ο (seealso ?2).. aud the response of the mass ratio to the change in donor mass We can then compute cep from refeq:zeta,/. usingthenecessaryparts from οστ Ον. θα μα»..."," The second term in \\ref{eq:zeta_rl} consists of the Roche lobe's response to the change in mass ratio, which can be described using Eggleton's approximation \citep{Egg83}: \citep[see also][]{SPV97}, , and the response of the mass ratio to the change in donor mass We can then compute $\zeta_\mathrm{RL}$ from \\ref{eq:zeta_rl}, using the necessary parts from \\ref{eq:dlna_dlnm}, \ref{eq:dlnq_dlnm} and \ref{eq:dlnrl_dlnq}."
+ By comparing Gaq and Crete)j we can find a μασ.Mae}. such that for all 3μαι Qaaμι) aud therefore AIT will be dwuamically stable.," By comparing $\zeta_\mathrm{ad}$ and $\zeta_\mathrm{RL}(\beta)$ we can find a $\beta_\mathrm{max}(q,m_\mathrm{d,c})$, such that for all $\beta\le \beta_\mathrm{max}$, $\zeta_\mathrm{ad}\ga \zeta_\mathrm{RL}(\beta)$ and therefore MT will be dynamically stable."
+ WeX visualize this in refzotas— for the case of a L2. red giaut. with uuchangies mass during its evolution. aud a 1.1M. companion at the start of MT. demonstrating how this condition changes for increasingly evolved red eiauts.," We visualize this in \\ref{zetas} for the case of a $1.2\,M_\odot$ red giant, with unchanging mass during its evolution, and a $1.1\,M_\odot$ companion at the start of MT, demonstrating how this condition changes for increasingly evolved red giants."
+ It can be seen from refzetas that in all the cases of fully nou-conservative MT (9= 0). the AIT is dynamically stable.," It can be seen from \\ref{zetas} that in all the cases of fully non-conservative MT $\beta=0$ ), the MT is dynamically stable."
+ We also note that fully conservative MT is dyuiuuically stable for cores with masses msm0.03AF... and in this case AIT will proceed ou the thermal timescale.," We also note that fully conservative MT is dynamically stable for cores with masses $m_\mathrm{d,c} \ga 0.63\,M_\odot$, and in this case MT will proceed on the thermal timescale."
+ Iu owe see that iu the more realistic case in which we account for mass loss (includiug wind loss). as well as z = 0.05. our ranee ofstalülitv is extended even further. to the xot of allowing tally conservative ATT for large core masses (λος0.16A. ).," In \\ref{realistic} we see that in the more realistic case in which we account for mass loss (including wind loss), as well as z = 0.03, our range ofstability is extended even further, to the point of allowing fully conservative MT for large core masses $M_\mathrm{d,c} \ga 0.46\,M_\odot$ )."
+ With decreasing, With decreasing
+transitions of aand'*CO.. at the position of the ppeak. using the IRAM 30mm telescope on 21 August 2010.,"transitions of and, at the position of the peak, using the IRAM m telescope on 21 August 2010."
+ These observations used the backend VESPA., These observations used the backend VESPA.
+ The spectra were smoothed to a velocity resolution of 1 ffor all the CO lines., The spectra were smoothed to a velocity resolution of 1 for all the CO lines.
+ The forward and beam efficiencies are and respectively for the (1—0) transition., The forward and beam efficiencies are and respectively for the (1–0) transition.
+ The same quantities are and for the (2-1) transition., The same quantities are and for the (2–1) transition.
+ The power beam widths for the (1-0) and (2-1) transitions are aand12.. respectively.," The half-power beam widths for the (1–0) and (2–1) transitions are and, respectively."
+ We compare the Herschel data of the 3302 region. with the VVLA and CO(2-1) HERA/30mmap at 12” resolution presented by Gratieretal.(2010).," We compare the Herschel data of the 302 region, with the VLA and CO(2–1) HERA/30mmap at $12''$ resolution presented by \citet{gratier2010}."
+.. We refer to the latter paper for a presentation of the noise properties., We refer to the latter paper for a presentation of the noise properties.
+ We also use theSpitzer MIPS mmap presented by Tabatabaetetal.(2007).. the KPNO mmap (Hoopes&Walterbos2000).. and the PACS and SPIRE maps at 100. 160. 250. 350 and citepkramer2010..verley2010.. boquien2010..," We also use the MIPS map presented by \citet{tabatabaei2007}, the KPNO map \citep{hoopes2000}, and the PACS and SPIRE maps at 100, 160, 250, 350 and \\citep{kramer2010, verley2010, boquien2010}."
+ The angular resolution of the 100 and 160 PPACS maps are ~6” aandΕ, The angular resolution of the 100 and 160 PACS maps are $\sim 6$ and.
+"λ, The rms noise levels of the PACS maps are 2.6mmly pix at mi aand 6.9mmJy pix!∣ at 160m wwhere the pixel sizes are 3722 and 6744 respectively."," The rms noise levels of the PACS maps are mJy $^{-2}$ at $\,\mu$ and mJy $^{-1}$ at $\,\mu$ where the pixel sizes are 2 and 4 respectively."
+ oobservations at two positions within our mapped region were extracted from the ISO/LWS archival data., observations at two positions within our mapped region were extracted from the ISO/LWS archival data.
+" The two positions are at RA= 01 3407 (42000) and RA=01""34""09° ((J2000).", The two positions are at RA= $^{\rm h}$ $^{\rm m}$ $^{\rm s}$ (J2000) and $^{\rm h}$ $^{\rm m}$ $^{\rm s}$ (J2000).
+ Within the 2’x region mapped with PACS reffig Hiresult)). wehavedetectedextended|C eemission Fromtayhenorthernspiralarmtracede.g.bythe 100 jm emission. with the strongest emission arising towards the rregion 3302 and (b) from the diffuse regions to the south-east and north-west.," Within the $2'\times2'$ region mapped with PACS \\ref{fig_ciiresult}) ), we have detected extended emission from (a) the northern spiral arm traced e.g. by the $100\,\mu$ m emission, with the strongest emission arising towards the region 302 and (b) from the diffuse regions to the south-east and north-west."
+ Comparison of the PPACS intensities with the ISO/LWS ddata at the two positions shows an agreement of better than at both positions., Comparison of the PACS intensities with the ISO/LWS data at the two positions shows an agreement of better than at both positions.
+ For the comparison of the intensities at the LWS positions. we first convolved the PACS mmap to the angular resolution of the LWS data.," For the comparison of the intensities at the LWS positions, we first convolved the PACS map to the angular resolution of the LWS data."
+ The full ISO/LWS ddata set along the major axis of M33 will be published in a separate paper by Abreu et al. (, The full ISO/LWS data set along the major axis of M33 will be published in a separate paper by Abreu et al. (
+1n prep.).,in prep.).
+ Some of the PACS 4m sspectra displayed baseline problems. attributed to the now decommissioned wavelength switching mode.," Some of the PACS $\,\mu$ spectra displayed baseline problems, attributed to the now decommissioned wavelength switching mode."
+ Spectra taken along the eemitting part of the spiral arm extending north-east to south- showed problems and have been blanked., Spectra taken along the emitting part of the spiral arm extending north-east to south-west showed problems and have been blanked.
+ Both. the um aand the um lines were detected at only a few positions within the mapped region.," Both, the $\,\mu$ and the $\,\mu$ lines were detected at only a few positions within the mapped region."
+" The maps of um aand um eemission towards the rregion look very similar and both peak at RA = 01 34"" 06:33. Dec = 30*447'223""7330."," The maps of $\,\mu$ and $\,\mu$ emission towards the region look very similar and both peak at RA = $^{\rm h}$ $^{\rm m}$ 3, Dec = 30."
+" In addition. the um mmap shows a secondary peak towards the south-west. to the south of the ridge. at RAZ01334""00:53364. Dee = 3074462267555 (J2000). which is not found in the mmap."," In addition, the $\,\mu$ map shows a secondary peak towards the south-west, to the south of the ridge, at 364 Dec = 55 (J2000), which is not found in the map."
+ The second ppeak lies between the two ridges detected in CO(2-1) and coincides with an ppeak., The second peak lies between the two ridges detected in CO(2–1) and coincides with an peak.
+ This suggests that the eemission at this peak position arises in very dense atomic gas., This suggests that the emission at this peak position arises in very dense atomic gas.
+ Overlays of the mmap with maps ofHa..Hi. CO(2-I) and dust continuum in the MIR and FIR (MIPS jim. PACS 100 um) are shown in Fig. 5..," Overlays of the map with maps of, CO(2–1) and dust continuum in the MIR and FIR (MIPS $\,\mu$ m, PACS 100 $\,\mu$ m) are shown in Fig. \ref{fig_allmaps}."
+ The dust continuum maps correlate well with the mmap., The dust continuum maps correlate well with the map.
+ They peak towards the rregion and show the spiral arm extending from the rregion in south-western direction., They peak towards the region and show the spiral arm extending from the region in south-western direction.
+ In contrast. CO emission shows a clumpy structure wrapping around the rregion towards the east.," In contrast, CO emission shows a clumpy structure wrapping around the region towards the east."
+ CO emission shows the spiral arm seen in the continuum. but its peaks are shifted towards the south.," CO emission shows the spiral arm seen in the continuum, but its peaks are shifted towards the south."
+ The eemission shows a completely different morphology. peaking towards the north and south of the rregion and showing a clumpy filament running towards the west.," The emission shows a completely different morphology, peaking towards the north and south of the region and showing a clumpy filament running towards the west."
+ Further below. we will discuss the correlations 1n more detail.," Further below, we will discuss the correlations in more detail."
+ Fig., Fig.
+ 6 shows overlays of the uim mmap at a resolution of wwith aand CO(2-1).," \ref{fig_o63olay} shows overlays of the $\,\mu$ map at a resolution of with and CO(2–1)."
+ Towards the south and south-west of the rregion. the um eemission matches the eemission well.," Towards the south and south-west of the region, the $\,\mu$ emission matches the emission well."
+ This is surprising given the high excitation requirements for the lline., This is surprising given the high excitation requirements for the line.
+ The secondary um ppeak towards the south-west. is also traced by Ητ.," The secondary $\,\mu$ peak towards the south-west, is also traced by ."
+. It lies in between two ridges of CO emission. the arm running north-east to south-west. and a second ridge of emission running in north-south direction.," It lies in between two ridges of CO emission, the arm running north-east to south-west, and a second ridge of emission running in north-south direction."
+ The northern part of this second ridge shows an interesting layering of emission: both, The northern part of this second ridge shows an interesting layering of emission: both
+3.3 Results from analysis,$\Sigma^-$ ) in Fig.9 as results from this analysis.
+," For comparison, we show the CTEQ3 input"
+slightly higher than the curve predicted: by the models in he brightest magnitude bins.,slightly higher than the curve predicted by the models in the brightest magnitude bins.
+ Since the data show evidence for a significant. [raction of binary stars among the cluster members. we have taken hem into account in the synthetic diagram.," Since the data show evidence for a significant fraction of binary stars among the cluster members, we have taken them into account in the synthetic diagram."
+ A mass ratio has »een associated to each system via random extractions from. a [lat cüstribution., A mass ratio has been associated to each system via random extractions from a flat distribution.
+ We have then followed the prescriptions eiven by Maeder (1974) to attribute colours and magnitudes ο svstenis with dillerent primary/secondary mass ratios (see Bragaglia ct al., We have then followed the prescriptions given by Maeder (1974) to attribute colours and magnitudes to systems with different primary/secondary mass ratios (see Bragaglia et al.
+ LOOT for more details)., 1997 for more details).
+ Ht is interesting to notice that all the stellar models lead to synthetic CMDSs in »etter agreement with the data when the assumed fraction of Xnaries is 304., It is interesting to notice that all the stellar models lead to synthetic CMDs in better agreement with the data when the assumed fraction of binaries is $\%$.
+ As already discussed by Maeder (1974) ancl Fan et al. (, As already discussed by Maeder (1974) and Fan et al. (
+1996). this percentage is larger that that derived ον counting the stars on the right side of the MS (see section 3.3) and the dillerence is due to the combination of several selection effects.,"1996), this percentage is larger that that derived by counting the stars on the right side of the MS (see section 3.3) and the difference is due to the combination of several selection effects."
+ We wish to note here that aff of the simulated: cases vield. rec-eiant branches that are slightly redeer than he observed ones and that in all cases. except in the FRANEC models with Z=0.02. this colour excess is already visible at the base of the RGB.," We wish to note here that $all$ of the simulated cases yield red-giant branches that are slightly redder than the observed ones and that in all cases, except in the FRANEC models with Z=0.02, this colour excess is already visible at the base of the RGB."
+ We ascribe this cllect to he uncertainty in the photometric conversions from. the heoretical plane (luminosity and effective. temperature) o the observational one (magnitude and colour): the emperature-colour conversion. is the one that djs most allected., We ascribe this effect to the uncertainty in the photometric conversions from the theoretical plane (luminosity and effective temperature) to the observational one (magnitude and colour): the temperature-colour conversion is the one that is most affected.
+" The fact that. at the same colours. the discrepancy is not observed. in the lower MS. (except. perhaps in. the ""acdova models with Z=0.02) suggests that the inadedquaevy of the temperature-colour. conversion does not concern all he cool stars. but. only those in low-eravity conditions. a circumstance foreseen by Bessel et al. ("," The fact that, at the same colours, the discrepancy is not observed in the lower MS (except perhaps in the Padova models with Z=0.02) suggests that the inadequacy of the temperature-colour conversion does not concern all the cool stars, but only those in low-gravity conditions, a circumstance foreseen by Bessel et al. ("
+1989).,1989).
+ Both of the examined sets of FRANEC evolutionary tracks provide svnthetic CMDs and Les in agreement with the observed data. although with slight dillerences in the choices of the parameters.," Both of the examined sets of FRANEC evolutionary tracks provide synthetic CMDs and LFs in agreement with the observed data, although with slight differences in the choices of the parameters."
+ With the set at Z=0.01. the best. fit is obtained. assuming an age of 1.6 Gyr. à reddening V)200.07 and a distance modulus (m-M)o27112.6.," With the set at Z=0.01, the best fit is obtained assuming an age of 1.6 Gyr, a reddening 0.07 and a distance modulus 12.6."
+ As shown in panels (a) and (d) of Fig., As shown in panels (a) and (d) of Fig.
+ 7. these stellar models reproduce quite well all the observed. features of the cluster (namely: the magnitude and colour distribution of the stars. the relative number of stars in the various evolutionary phases and their morphology. including the MS gaps).," \ref{fig-sim} these stellar models reproduce quite well all the observed features of the cluster (namely: the magnitude and colour distribution of the stars, the relative number of stars in the various evolutionary phases and their morphology, including the MS gaps)."
+ A very good agreement with the data. practically uncdistinguishable from that shown in Fig.," A very good agreement with the data, practically undistinguishable from that shown in Fig."
+ 7. (a) and (d). is achieved also with the FRANEC set with Z=0.02.," \ref{fig-sim} (a) and (d), is achieved also with the FRANEC set with Z=0.02."
+ In this case. the best model indicates roughly the same age { Gyr) ancl distance modulus ((m-M)o—112.7). but a loweες reddening V)=00.01). to compensate the intrinsically redder colours. corresponding to the doubled: metallicity.," In this case, the best model indicates roughly the same age (1.5 Gyr) and distance modulus 12.7), but a lower reddening 0.01), to compensate the intrinsically redder colours corresponding to the doubled metallicity."
+- These tracks have a metal content nominally larger than that spectroscopically attributed to NGC2506., These tracks have a metal content nominally larger than that spectroscopically attributed to NGC2506.
+ Llowever. according to their authors. the FRANEC tracks (i.0.. both the set with Z=0.01 and with Z=0.02) actually correspond to models with metallicity about hall of their nominal value. once the clleet of using the old. LAOL opacitics rather than the most recent OPAL opacitics is taken into account.," However, according to their authors, the FRANEC tracks (i.e., both the set with Z=0.01 and with Z=0.02) actually correspond to models with metallicity about half of their nominal value, once the effect of using the old LAOL opacities rather than the most recent OPAL opacities is taken into account."
+ Both sets are. therefore. roughly compatible with the spectroscopic metallicity 0.5 derived by Friel Janes (1993).," Both sets are, therefore, roughly compatible with the spectroscopic metallicity $-$ 0.5 derived by Friel Janes (1993)."
+99% levels is much greater than the range reported. by. 2..,$99$ levels is much greater than the range reported by \citet{Edelmann05}.
+ We therefore cannot rule out small eccentricities comparable to those reported by 2? in any of these binaries., We therefore cannot rule out small eccentricities comparable to those reported by \citet{Edelmann05} in any of these binaries.
+ 1n Table 1 we list our racial velocity measurements for the 108 remaining sdDs which we have observed as a part of this project. but not previously. published.," In Table \ref{tab:rv2a} we list our radial velocity measurements for the 108 remaining sdBs which we have observed as a part of this project, but not previously published."
+ La ss of these svstenis. we detect no significant radial velocity. variation.," In 88 of these systems, we detect no significant radial velocity variation."
+ ‘These systems are likely to be either single sdBs or binary systems in which the mass function is too low for us to detect radial velocity. variations., These systems are likely to be either single sdBs or binary systems in which the mass function is too low for us to detect radial velocity variations.
+ The remaining 20 sdBs do show significant racial velocity variations. and we consider it likely that most or all of these are binaries.," The remaining 20 sdBs do show significant radial velocity variations, and we consider it likely that most or all of these are binaries."
+ Our current data are not sullicient to distinguish the true orbital period in these systems from. various competing aliases. and so these targets are prime candidates for future measurements.," Our current data are not sufficient to distinguish the true orbital period in these systems from various competing aliases, and so these targets are prime candidates for future measurements."
+ We mark these systems in Fable 1 with an asterisk., We mark these systems in Table \ref{tab:rv2a} with an asterisk.
+ Two systems of particular note are PGI610|519 and. PG23171046., Two systems of particular note are PG1610+519 and PG2317+046.
+ Phe data we have collected. to date suggest that the orbital periods of these systems may be relatively long at 50° 70 davs., The data we have collected to date suggest that the orbital periods of these systems may be relatively long at $50$ – $70$ days.
+ μις list of twenty candidate binaries is not exhaustive., This list of twenty candidate binaries is not exhaustive.
+ As we have previouslv. discussed. some fraction of the," As we have previously discussed, some fraction of the"
+the most extended clusters in the outer Galactic halo and nearby dwarf galaxies.,the most extended clusters in the outer Galactic halo and nearby dwarf galaxies.
+ It would be important to know if any of the largest faint fuzzies lie above our Eq. 1..," It would be important to know if any of the largest faint fuzzies lie above our Eq. \ref{e:cutoff},"
+ as appears possible from Figure 7.., as appears possible from Figure \ref{f:rhvsmvexternal}.
+ Presently available observations appear to be consistent with the iypothesis that the globular clusters in the outer (Av...15 Kpe) Galactic halo were once all associated with dwarf spheroidal-like Tagments that have since disintegrated., Presently available observations appear to be consistent with the hypothesis that the globular clusters in the outer $R_{\rm{gc}} > 15$ kpc) Galactic halo were once all associated with dwarf spheroidal-like fragments that have since disintegrated.
+" On the other hand the majority of inner halo globular clusters with //,.«LO Kpe were orobably formed in association with the main body of the Milky Way system.", On the other hand the majority of inner halo globular clusters with $R_{\rm{gc}} < 10$ kpc were probably formed in association with the main body of the Milky Way system.
+ The fact that the sample of outer halo globulars contains more faint clusters than the inner halo is likely due to he destruction of low-mass inner clusters by disk shocks and idal stripping., The fact that the sample of outer halo globulars contains more faint clusters than the inner halo is likely due to the destruction of low-mass inner clusters by disk shocks and tidal stripping.
+ By the same token the scarcity of large globulars at small #4. values is likely also due to such destructive forces., By the same token the scarcity of large globulars at small $R_{\rm{gc}}$ values is likely also due to such destructive forces.
+ The presence of a few quite metal-rich clusters. such as Terzan 7 CFe/H]= 0.58) and Palomar | CFe/H]= 0.60). and perhaps Palomar 12 Fe/H]= 0.98) at quite large Galactocentric radii appears anomalous.," The presence of a few quite metal-rich clusters, such as Terzan 7 $[$ $/$ $] = -0.58$ ) and Palomar 1 $[$ $/$ $] = -0.60$ ), and perhaps Palomar 12 $[$ $/$ $] = -0.98$ ) at quite large Galactocentric radii appears anomalous."
+ The existence of such metal-rich objects in the outer Galactic halo can be explained if they formed in. or in association with. dwarf spheroidal galaxies (as appears likely for Ter.," The existence of such metal-rich objects in the outer Galactic halo can be explained if they formed in, or in association with, dwarf spheroidal galaxies (as appears likely for Ter."
+ 7 and Pal., 7 and Pal.
+ 12)., 12).
+ With one exception. luminous globular clusters in the outer halo are all compact whereas faint ones may have any radius. a result which also holds for globular clusters in the LMC. SMC and Fornax dwarf spheroidal.," With one exception, luminous globular clusters in the outer halo are all compact whereas faint ones may have any radius, a result which also holds for globular clusters in the LMC, SMC and Fornax dwarf spheroidal."
+" We speculate that the luminous (Aly= 9.58) and very large (22),=17.9 po) cluster NGC 2419. which is located at a Galactocentric distance of 91.5 Kpe. might be the remnant core of a now dispersed dwarf spheroidal galaxy."," We speculate that the luminous $M_V = -9.58$ ) and very large $R_h = 17.9$ pc) cluster NGC 2419, which is located at a Galactocentric distance of $91.5$ kpc, might be the remnant core of a now dispersed dwarf spheroidal galaxy."
+ Apart from apparently not possessing an internal metallicity spread. its properties are similar to the very large and luminous globular cluster :; Centauri. which might also be such a stripped core of a former dwarf spheroidal galaxy.," Apart from apparently not possessing an internal metallicity spread, its properties are similar to the very large and luminous globular cluster $\omega$ Centauri, which might also be such a stripped core of a former dwarf spheroidal galaxy."
+ We thank Paul Hodge for reminding us of his prescient 1962 paper on the diameter and structure of clusters in the LMC., We thank Paul Hodge for reminding us of his prescient 1962 paper on the diameter and structure of clusters in the LMC.
+ ADM recognises financial support from PPARC in the form of a Postdoctoral Fellowship. and is grateful to Mark Wilkinson for a useful discussion.," ADM recognises financial support from PPARC in the form of a Postdoctoral Fellowship, and is grateful to Mark Wilkinson for a useful discussion."
+"work of Shaviv&Starrfield(1987).. and about half the estimated value of who used azz0.2L,,,.","work of \citet{sha87}, and about half the estimated value of \citet{reg89} who used $\alpha \approx 0.2 L_{acc}$."
+" We then asstune different rotation rates For a non rotating star 7)—0 and the BL luminosity is exactly half the accretion energy Ly,=fLoe2. while lor a star rotating near break-up 7j—1 and Ly,=0."," We then assume different rotation rates For a non rotating star $\eta = 0$ and the BL luminosity is exactly half the accretion energy $L_{BL}=L_{acc}/2$, while for a star rotating near break-up $\eta=1$ and $L_{BL}=0$."
+ In this work we took values ranging from η=| (for no boundary laver irradiation. in order to check only the effect of the compressional heating) down to 7=0.05 (when the star is slowly rotating and boundary laver irradiation is (he main source of heating).," In this work we took values ranging from $\eta = 1$ (for no boundary layer irradiation, in order to check only the effect of the compressional heating) down to $\eta = 0.05$ (when the star is slowly rotating and boundary layer irradiation is the main source of heating)."
+ Clearly a smaller value of a will require a smaller value of jj in order to keep the same amount of DL irradiation in a specific model. ancl vice versa.," Clearly a smaller value of $\alpha$ will require a smaller value of $\eta$ in order to keep the same amount of BL irradiation in a specific model, and vice versa."
+ However. (he compressional heating results obtained in this work are not at all affected bv the value of à. used in the simulations.," However, the compressional heating results obtained in this work are not at all affected by the value of $\alpha$ used in the simulations."
+ In (ihe simulations it is assumed (hat the accretion and heating of the white dwarf is uniform rather than being restricted to the equatorial region., In the simulations it is assumed that the accretion and heating of the white dwarf is uniform rather than being restricted to the equatorial region.
+ In addition. the transfer of angular momentum (by shear mixing) into the white dwarl is neglected.," In addition, the transfer of angular momentum (by shear mixing) into the white dwarf is neglected."
+" As previously remarked in section 2. during the outburst. as accretion takes place at a high rate. (he stars photospheric emission is overwhelmed by (he emission of the hot components (mainly the inner disk). which makes it difficult to assess the exact temperature ol the star and its rotation rate €,"," As previously remarked in section 2, during the outburst, as accretion takes place at a high rate, the star's photospheric emission is overwhelmed by the emission of the hot components (mainly the inner disk), which makes it difficult to assess the exact temperature of the star and its rotation rate $\Omega_{wd}$."
+ ILowever. on day 53 the svstem is found in a low state and starts to cool.," However, on day 53 the system is found in a low state and starts to cool."
+ During that Gime the accretion rate has probably dropped toits quiescence level., During that time the accretion rate has probably dropped to its quiescence level.
+ Numerical simulations (Godon&Sion2002) have shown that the temperature increase. due to BL irradiation. is sustained only curing accretion. and when the accretion is turned olf. the star rapidly radiates away the DL energy absorbed in its outermost laver.," Numerical simulations \citep{god02} have shown that the temperature increase, due to BL irradiation, is sustained only during accretion, and when the accretion is turned off, the star rapidly radiates away the BL energy absorbed in its outermost layer."
+ However. (he temperature increase due (o compressional heating takes place deeper in the outer lavers ol the star and it takes many days (months) for the star to cool.," However, the temperature increase due to compressional heating takes place deeper in the outer layers of the star and it takes many days (months) for the star to cool."
+ Therefore. we assume that in the cooling phase the observed elevated temperature of the star is due to the compressional heatinge it has endured duringe the outburst phase alone.," Therefore, we assume that in the cooling phase the observed elevated temperature of the star is due to the compressional heating it has endured during the outburst phase alone."
+ since DL irradiation and compressional heating take place at different depths in (he outer lavers of the star. and on different Gme scales. we model their effect on the temperature of (he star in separate model runs.," Since BL irradiation and compressional heating take place at different depths in the outer layers of the star, and on different time scales, we model their effect on the temperature of the star in separate model runs."
+made using single temperature models fitted to gas with a mean temperature of 2-4 keV. If the gas has a complex. multi-temperature structure. abundance may be overestimated owing either to the mixture of differing levels of Fe-L and Fe-K shell emission along the line of sight (2).. or the behaviour of the spectral fitting code when a single temperature model is used to describe a temperature spectrum (2)..,"made using single temperature models fitted to gas with a mean temperature of 2-4 keV. If the gas has a complex, multi-temperature structure, abundance may be overestimated owing either to the mixture of differing levels of Fe-L and Fe-K shell emission along the line of sight \citep{Rasiaetal08}, or the behaviour of the spectral fitting code when a single temperature model is used to describe a multi-temperature spectrum \citep{Gastaldelloetal10}."
+ Testing for this type of problem is difficult. in that high signal-to-noise spectra are required if multi-temperature models are to be fitted successfully. but larger regions are more likely to contain significant temperature variations.," Testing for this type of problem is difficult, in that high signal–to–noise spectra are required if multi-temperature models are to be fitted successfully, but larger regions are more likely to contain significant temperature variations."
+ To determine whether the inverse Fe-bias is likely to affect our abundance measurements. we performed three tests: A powerlaw emission component associated with the radio jets or lobes could also affect the measured abundance.," To determine whether the inverse Fe-bias is likely to affect our abundance measurements, we performed three tests: A powerlaw emission component associated with the radio jets or lobes could also affect the measured abundance."
+ If the powerlaw slope is similar to that of the bremsstrahlung continuum. fitting the spectrum with a simple plasma model is likely to result in an underestimate of abundance.," If the powerlaw slope is similar to that of the bremsstrahlung continuum, fitting the spectrum with a simple plasma model is likely to result in an underestimate of abundance."
+ If the powerlaw provides more emission at higher energies. the plasma model temperature is likely to be biased high. and abundances overestimated.," If the powerlaw provides more emission at higher energies, the plasma model temperature is likely to be biased high, and abundances overestimated."
+ The expected flux from inverse Compton scattering in the radio lobes is too low to be detected. and indeed no evidence powerlaw emission is found (see paper D.," The expected flux from inverse Compton scattering in the radio lobes is too low to be detected, and indeed no evidence powerlaw emission is found (see paper I)."
+ To test whether there could be significant powerlaw emission from the jets. we tit apectpowerlaw models to the spectra extracted in the E-W and N-S profiles.," To test whether there could be significant powerlaw emission from the jets, we fit apec+powerlaw models to the spectra extracted in the E–W and N–S profiles."
+ If à powerlaw componen were associated with the jets we would expect to see the highes fluxes from regions 3-5 of the E-W protile. and none in regions I. 4 and 5 of the N-S profile.," If a powerlaw component were associated with the jets we would expect to see the highest fluxes from regions 3-5 of the E–W profile, and none in regions 1, 4 and 5 of the N–S profile."
+ In practise we find powerlaw contributions o be consistent with zero in all regions except regions 2 and 3 of he E-W profile. and region | of the N-S profile.," In practise we find powerlaw contributions to be consistent with zero in all regions except regions 2 and 3 of the E–W profile, and region 1 of the N–S profile."
+ In these regions he best-fitting powerlaw index is inverted. and the component only oroduees significant additional flux above 5 keV. If the powerlaw index is fixed at Γ 5-2. as expected from the radio spectral index (see προς D. the powerlaw flux is consistent with zero.," In these regions the best-fitting powerlaw index is inverted, and the component only produces significant additional flux above 5 keV. If the powerlaw index is fixed at $\Gamma$ =2, as expected from the radio spectral index (see paper I), the powerlaw flux is consistent with zero."
+ We therefore conclude that no significant powerlaw emission is detected from he jet. and that the abundances are not biased in this way.," We therefore conclude that no significant powerlaw emission is detected from the jet, and that the abundances are not biased in this way."
+ High abundances could also be found in error if the plasma is not in ionisation equilibrium., High abundances could also be found in error if the plasma is not in ionisation equilibrium.
+ ο point out that the Maxwellian electron distribution assumed for thermal plasmas is not valid in shocked regions., \citet{Kaastraetal09} point out that the Maxwellian electron distribution assumed for thermal plasmas is not valid in shocked regions.
+ Non-thermal electrons alter the X-ray spectrum and ean lead to the overestimation of the abundance: À possible example is seen in the high abundance are associated with the shock in HCG 62 ¢?).., Non-thermal electrons alter the X–ray spectrum and can lead to the overestimation of the abundance; A possible example is seen in the high abundance arc associated with the shock in HCG 62 \citep{Gittietal10}.
+ Both shocks or mixing of non-thermal electrons into the thermal plasma could occur along the edge of the radio Jets in AWM 4., Both shocks or mixing of non-thermal electrons into the thermal plasma could occur along the edge of the radio jets in AWM 4.
+ The weak surface brightness features associated with the jets could indicate the presence of compressed gas (see paper I. fig.," The weak surface brightness features associated with the jets could indicate the presence of compressed gas (see paper I, fig."
+ 2)., 2).
+ However. we might expect that shocks capable of affecting the abundance measurements would be detectable as either density or temperature features.," However, we might expect that shocks capable of affecting the abundance measurements would be detectable as either density or temperature features."
+ The only regions where such features may be visible are in the western knot and eastern bend of the jets. but our data are insuftficiently deep to confirm or reject this.," The only regions where such features may be visible are in the western knot and eastern bend of the jets, but our data are insufficiently deep to confirm or reject this."
+ We also note that we might expect to find both mixing of relativistic electrons from the radio source and multi-phase gas in the lobes of the radio source as well as the jets. since the lobes appear to contain a mix of thermal and relativistic plasmas (see paper D.," We also note that we might expect to find both mixing of relativistic electrons from the radio source and multi-phase gas in the lobes of the radio source as well as the jets, since the lobes appear to contain a mix of thermal and relativistic plasmas (see paper I)."
+ The lack of high abundances coincident with the lobes therefore argues against these sources of bias being effective., The lack of high abundances coincident with the lobes therefore argues against these sources of bias being effective.
+ The small extension of the high abundance region north of the galaxy core also cannot be explained as a shock or jet-related feature., The small extension of the high abundance region north of the galaxy core also cannot be explained as a shock or jet-related feature.
+ We therefore conclude that while it is likely that the spectra used in the maps do contain emission from multiple plasma components with different temperatures and abundances. the variation is probably not large enough to affect our results. and the single-temperature fits provide a sufficiently accurate estimate of abundance for our purposes.," We therefore conclude that while it is likely that the spectra used in the maps do contain emission from multiple plasma components with different temperatures and abundances, the variation is probably not large enough to affect our results, and the single-temperature fits provide a sufficiently accurate estimate of abundance for our purposes."
+"(for LO modes. and 6£/£ = 0.06. 814,;, decreases from 2.| plz ou the coarseerid to 0.5 gllz on the fine grid) or the radius difference (for 10 amodes. and 6R/R = 0.03. 14,4 decreases frou 6.1 μΗΣ ou the coarse grid to 1.0 plz ou the fine grid).","(for 40 modes, and $ \delta L / L$ = 0.06, $ \delta \nu_{min}$ decreases from 2.4 $\mu$ Hz on the coarsegrid to 0.5 $\mu$ Hz on the fine grid) or the radius difference (for 40 modes, and $ \delta R / R$ = 0.03, $ \delta \nu_{min}$ decreases from 6.4 $\mu$ Hz on the coarse grid to 1.0 $\mu$ Hz on the fine grid)."
+" Since we cannot freely refine the eric because of computing time. the values obtained for 67,5), ou the refined erid may be overestimated: we caunot exclude that a finer eid. gives sumaller values of ορ."," Since we cannot freely refine the grid because of computing time, the values obtained for $ \delta \nu_{min}$ on the refined grid may be overestimated: we cannot exclude that a finer grid gives smaller values of $ \delta \nu_{min}$."
+ AMauy asteroscisinological studies do not directly compare he pulsation frequencies of two stars but rather compare he large aud simall separations., Many asteroseismological studies do not directly compare the pulsation frequencies of two stars but rather compare the large and small separations.
+ For this secoud method. we define the difference 944 between two spectra as the uaxiuun value of the large separation or of the siall separation (the highest value is retained).," For this second method, we define the difference $ \delta_{sls}$ between two spectra as the maximum value of the large separation or of the small separation (the highest value is retained)."
+" The results for he coarse erid. for the mass difference óAM/A with 60 considered modes. are shown in Fie. ὃν,"," The results for the coarse grid, for the mass difference $ \delta M / M$ with 60 considered modes, are shown in Fig. \ref{petitesepa}."
+ daz. is about a Actor of two lower hau à»., $ \delta_{sls}$ is about a factor of two lower than $\delta \nu $.
+" For the luuinosity difference. us Is also about a factor of two smaller than ὃν, For the radius clifference. 84, Is about a factor of three lower tha o."," For the luminosity difference, $ \delta_{sls}$ is also about a factor of two smaller than $\delta \nu $ For the radius difference, $ \delta_{sls}$ is about a factor of three lower than$\delta \nu $ ."
+ As in paragraph 3.2.. the same behaviow is found when using the fiue exid fora given 0M/ AL. 8L/Lor ORE R.," As in paragraph \ref{diffreq}, , the same behaviour is found when using the fine grid for a given $ \delta M / M$ , $ \delta L / L$or $ \delta R / R$ ."
+"where à=910,h712.6 eV (ChanaudChu Q, and Ὁ are cosinological density parameters of neutrinos aud total matter. and P2:0.7 is the ITubble parameter.","where $\alpha=94 \Omega_mh^2 \approx 12.6$ eV \citep{Chan}, $\Omega_{\nu}$ and $\Omega_m$ are cosmological density parameters of neutrinos and total matter, and $h \approx 0.7$ is the Hubble parameter."
+ By combining Eqs. (, By combining Eqs. (
+8)-(12). we finally ect: From Eq. (,"8)-(12), we finally get: From Eq. ("
+"13). we notice that for fixedmn. rn, depends on 49 ouly.","13), we notice that for fixed$m_{\nu}$, $r_c$ depends on $u_0$ only."
+ Plotting 7? against kc. we see that the values of R are uearly coustaut for all clusters.," Plotting $R$ against $r_c$, we see that the values of $R$ are nearly constant for all clusters."
+" P is approximately proportional to 1/0, (see Fig.", $R$ is approximately proportional to $1/m_{\nu}$ (see Fig.
+ 2)., 2).
+" If my<2 eV. which is the current upper bound (Elearov 2007).. then Rrn, for most clusters."," If $m_{\nu} \le 2$ eV, which is the current upper bound \citep{Oystein,Sanders}, then $R \gg r_c$ for most clusters."
+ Suppose the total central denusitv of a cluster py is related to p. by a power law p.~py: by defining lar?py=dM(r)fdr at r=0 aud rearranging Eq. (," Suppose the total central density of a cluster $\rho_0$ is related to $\rho_c$ by a power law $\rho_c \sim \rho_0^{\gamma}$; by defining $4 
+\pi r^2 \rho_0=dM(r)/dr$ at $r=0$ and rearranging Eq. ("
+8). we obtain the kev relatiouship between the core radius rand the product 2T: where we have assumed that αιας.|RePντ)Di ds+ rearly a constat for all clusters.,"8), we obtain the key relationship between the core radius $r_c$ and the product $\beta 
+T$: where we have assumed that $\ln \ln(1+R^2/r_c^2)$ is nearly a constant for all clusters."
+" To verify the above xediction. we plot Iup, against IntT) for 103 clusters in Fig.3: an approximately linear relation is obtained which aerees with Eq. ("," To verify the above prediction, we plot $\ln r_c$ against $\ln (\beta T)$ for 103 clusters in Fig.3; an approximately linear relation is obtained which agrees with Eq. ("
+11).,14).
+ The slope in Fig., The slope in Fig.
+" 25 is 1,97£0.11 which corresponds to 5%3/2 (correlation cocficient = 0.66).", 3 is $0.97 \pm 0.11$ which corresponds to $\gamma \approx -3/2$ (correlation coefficient $=0.66$ ).
+" ILowever. the wneertaintics in AL... n TF and +, are quite large. aud the total mass xofile of a cluster (Eq. ("," However, the uncertainties in $M_c$, $\beta$, $T$ and $r_c$ are quite large, and the total mass profile of a cluster (Eq. ("
+7)) is only derived. by using hiuss anuodel,7)) is only derived by using King's $\beta$ -model.
+" Therefore our model can only give au approximate prediction of the relation between lur, and InfoP) with >53/2."," Therefore our model can only give an approximate prediction of the relation between $\ln 
+r_c$ and $\ln (\beta T)$ with $\gamma \approx -3/2$."
+ Neutrinos exist in clusters aud they may fori structures with help of cold dark matter (ClanaudChu2006).., Neutrinos exist in clusters and they may form structures with help of cold dark matter \citep{formation}.
+" By assunüug the bydrostatic equilibrium of neutrinos and hot eas particles with total mass in clusters. we obtain the density profile of ucutrinos iu terms of 0, T and r2 and we can thereby obtain au approximate relation amoung these parameters with mj<2 eV. If peX Pye then a linear relationship between Iue aud InoF is obtained which agrees with the observed data with 5=— 23/2. Our result is also compatible with Sanders(2007) that the core profiles iu clusters can be explained by neutrinos as dark matter."," By assuming the hydrostatic equilibrium of neutrinos and hot gas particles with total mass in clusters, we obtain the density profile of neutrinos in terms of $\beta$, $T$ and $r_c$, and we can thereby obtain an approximate relation among these parameters with $m_{\nu} \le 2$ eV. If $\rho_c 
+\propto \rho_0^{\gamma}$ , then a linear relationship between $\ln r_c$ and $\ln \beta T$ is obtained which agrees with the observed data with $\gamma \approx -3/2$ Our result is also compatible with \citet{Sanders} that the core profiles in clusters can be explained by neutrinos as dark matter."
+ccould measure the temperature of any newly-discovered ccluster to with only a modest 50 ks exposure.,could measure the temperature of any newly-discovered cluster to with only a modest 50 ks exposure.
+ This opens the path to important science with large follow-up programs on aandC, This opens the path to important science with large follow-up programs on and.
+"handra.. A ~5 Ms program onXMM-Newton,, for instance, could measure profiles and temperatures for ~100cclusters at z>0.6, and masses for a large subset."," A $\sim 5$ Ms program on, for instance, could measure profiles and temperatures for $\sim 100$clusters at $z>0.6$, and masses for a large subset."
+" This would be a significant increase in sample size and lead to important advances in the use of clusters as cosmological probes, as well as in our understanding of cluster physics at higher redshifts."," This would be a significant increase in sample size and lead to important advances in the use of clusters as cosmological probes, as well as in our understanding of cluster physics at higher redshifts."
+" While ambitious, this type of program is feasible and falls naturally under the category of envisaged for after 2010."," While ambitious, this type of program is feasible and falls naturally under the category of envisaged for after 2010."
+" In conclusion, we expect the PCC to be unique for its ability to find rare, massive clusters out to high redshifts, and hence become a workhorse cluster catalog for many types of detailed cluster studies."," In conclusion, we expect the PCC to be unique for its ability to find rare, massive clusters out to high redshifts, and hence become a workhorse cluster catalog for many types of detailed cluster studies."
+" In particular, we expect the PCC to initiate important X-ray follow-up programs on aand Chandra.."," In particular, we expect the PCC to initiate important X-ray follow-up programs on and ."
+"with WIND is dimuner than that powered by the low c O1CS,",with WIND is dimmer than that powered by the low $\sigma$ ones.
+ It is well known that for the totally ordered magnetic configuration. high linear polarization is expected.," It is well known that for the totally ordered magnetic configuration, high linear polarization is expected."
+ For the random magnetic configuration. mild linear polarization can be expected if some specific ecometry effects have con. takeu iuto account.," For the random magnetic configuration, mild linear polarization can be expected if some specific geometry effects have been taken into account."
+ Sometimes the magnetic field is iot only ordered or onlv random., Sometimes the magnetic field is not only ordered or only random.
+ Tren. it is imteresting o investigate its near polarization., Then it is interesting to investigate its linear polarization.
+ Ifere we propose a siuple formula to describe the Linear polarization xoperties of a slab of such a mixed maguetic field. with which we can see the inpact of the ordered feld on the »oluization.," Here we propose a simple formula to describe the linear polarization properties of a slab of such a mixed magnetic field, with which we can see the impact of the ordered field on the polarization."
+" Following Laing (1050: See hiis Appeucdix Al for detail). he coordinates involved are defined as follows ( see figure 3): ais the anele between the plane of the ejecta aud the ine of sight: c. δν 2 are rectangular coordinates with the g-aNis poiutiug towar‘ds the observer (1ος, he «irection 11) and the y-axis paralcl to the “local” plaue of he ejecta: /εν a! are coordinates in the plane of t jecta. g ds parallel to y: 0 is the angle between the feld direction and the 4 axis at auy point in the ejecta: X is the position angle of the E-vector of the polarized radiation. measured from the wy plane."," Following Laing (1980; See his Appendix A1 for detail), the coordinates involved are defined as follows ( see figure 3): $\alpha$ is the angle between the plane of the ejecta and the line of sight; $x$, $y$, $z$ are rectangular coordinates with the z-axis pointing towards the observer (i.e., the direction $\textbf{n}$ ) and the y-axis parallel to the “local” plane of the ejecta; $x'$, $y'$ are coordinates in the plane of the ejecta, $y'$ is parallel to y; $\theta$ is the angle between the field direction and the $x'$ axis at any point in the ejecta; $\chi$ is the position angle of the E-vector of the polarized radiation, measured from the $x-y$ plane."
+" Therefore 1ο rancdoni (ordered) magueticteld vector Bray (Bora) at a poiut in the «lab By,=Που(ήsno.si1.cosÓà). BorgBo(cos(ysino.sindy.cosycosa) respectively."," Therefore the random (ordered) magnetic-field vector $\textbf{B}_{\rm ran}$ $\textbf{B}_{\rm
+ord}$ ) at a point in the slab $\textbf{B}_{\rm
+ran}=B(\cos \theta \sin \alpha,~\sin \theta,~\cos\theta \cos
+\alpha )$, $\textbf{B}_{\rm ord}=B_0(\cos \theta_0 \sin
+\alpha,~\sin \theta_0,~\cos\theta_0 \cos \alpha )$ respectively."
+ Thus the total maguetic field is where 5b=Bo/B., Thus the total magnetic field is where $b\equiv B_0/B$.
+ The clectric field of a linearly polarized clectromaguetic wave ids directed along the vector Now \ satisfies With equation (32). it is casy to get expressions for cos(2\) and sin(24).," The electric field of a linearly polarized electromagnetic wave is directed along the vector Now $\chi$ satisfies With equation (32), it is easy to get expressions for $\cos
+(2\chi)$ and $\sin(2\chi)$."
+" With equations (1)(3) of Laing (1980) and asstmineg the spectra index p=3. we have the Stokes parameters For 0,=7/2or32/2. we have ©=0 and the degree of the linear polarization For b<>1. H has the maxiuuu value II—i."," With equations (1)–(3) of Laing (1980) and assuming the spectra index $p=3$, we have the Stokes parameters For $\theta_0=\pi/2 ~{\rm or}~3\pi/2$, we have $U=0$ and the degree of the linear polarization For $b\gg1$, $\Pi$ has the maximum value $\Pi={3\over 4}$."
+ Iu the current work. &~0.05—l1. the correspouding toroidal uaenetie field is far stronger than that ogenerated iu reverse shock. 1.0... 59l. so the local poiut polarization can be as high as 75%.," In the current work, $\sigma\sim 0.05-1$, the corresponding toroidal magnetic field is far stronger than that generated in reverse shock, i.e., $b\gg1$, so the local point polarization can be as high as $75\%$."
+ For ultra-relativistic ο]ecta. due to he beaming effect. only. the euissiou coniιο fron) a very ight cone around the line of sight can beetected.," For ultra-relativistic ejecta, due to the beaming effect, only the emission coming from a very tight cone around the line of sight can be detected."
+ If the ine of sight is slightly offthe svaunetzic axis « the ordere naenetic field. t1ο orientaion of the viewed magnetic fiek is nearly the same.," If the line of sight is slightly off the symmetric axis of the ordered magnetic field, the orientation of the viewed magnetic field is nearly the same."
+ T1e hnieh lirear net polarization is expected since t16 local high linear polarization cannot be averaged effectively., The high linear net polarization is expected since the local high linear polarization cannot be averaged effectively.
+ The «etailed nuncerieal calculation of the net polarization wi] be preseated elsewhere., The detailed numerical calculation of the net polarization will be presented elsewhere.
+" Tere. for simplicity. following the treatment of Grano Isouniel (2003). the net Stokes parameters of the ordered magnetic field (Caq.Qua. fora} and of the random magnetic field (C.Qui. fran) are calculated separately,"," Here, for simplicity, following the treatment of Granot Könnigl (2003), the net Stokes parameters of the ordered magnetic field $U_{\rm ord},~Q_{\rm
+ord},~I_{\rm ord}$ ) and of the random magnetic field $U_{\rm ran},~Q_{\rm
+ran},~I_{\rm ran}$ ) are calculated separately."
+ Therefore, Therefore
+Optimal inversion techniques should avoid binning while relvine on techuiques such as Ἱνοπο interpolation.,Optimal inversion techniques should avoid binning while relying on techniques such as Kernel interpolation.
+ We ound here that for such a sample and when assessing quantities which are two derivatives away from the data. ie Wernel introduces spurious high frequency features in the recovered mass profile.," We found here that for such a sample and when assessing quantities which are two derivatives away from the data, the Kernel introduces spurious high frequency features in the recovered mass profile."
+ Binning the projected quantities on the other haud allows us to coutrol visually he quality of the fit., Binning the projected quantities on the other hand allows us to control visually the quality of the fit.
+ We use floating biuuiug which is defiuec as follows: for cach galaxw we find its p-nuearest reighbors. and define a ring which CLCOLIPASSCs thei exactly: the estimator for the density. Spine. would be defined as p divided by the area of that rine.," We use floating binning which is defined as follows: for each galaxy we find its p-nearest neighbors, and define a ring which encompasses them exactly; the estimator for the density, $\Sigma_{\R{Ring}}$, would be defined as p divided by the area of that ring."
+ For the srojected. velocityH dispersionBH squared. pD07. we could sum over the velocity squared (neasured with respect to the nean velocity of the cluster) of the p uciglibors aud divide *oprdn. practice. a better estimator.. 2a5. accounting. for. velocity iieasurement errors is given by where 97 is the stun of the error ou the measured variance. 07.2 and of the square of the measured error ou the velocity ¢;.," For the projected velocity dispersion squared, $\sigma _p^2$, we could sum over the velocity squared (measured with respect to the mean velocity of the cluster) of the p neighbors and divide by p; in practice a better estimator, $\sigma _p^{2\star}$, accounting for velocity measurement errors is given by where $\sigma _i^2$ is the sum of the error on the measured variance, $ \sigma _p^2$, and of the square of the measured error on the velocity $v_i$ ."
+ Bootstrapping is applied to estimate 0. while first ucelecting these measurement errors., Bootstrapping is applied to estimate $\sigma _p^2$ while first neglecting these measurement errors.
+ Aw estimate of the projected energv deusitv is even by unDespe ," An estimate of the projected energy density is given by $\sigma _p^{2\star}\ \times \Sigma
+_{\R{Ring}}$."
+In practice. binning over 10 to 15 neighboring ealaxies is applied. viclding estimates of the Poisson noise induced by sample.," In practice, binning over $10$ to $15$ neighboring galaxies is applied, yielding estimates of the Poisson noise induced by sampling."
+ The sample is truncated iu projected radius AR., The sample is truncated in projected radius $R$.
+ Since eenericallv. truncation aud deprojection will not commute. he estimation of the cumulative mass profile arising frou a truucated sample in projected radius will be biased.," Since generically, truncation and deprojection will not commute, the estimation of the cumulative mass profile arising from a truncated sample in projected radius will be biased."
+ Iu physical terms. this follows because we cannot distinguish vetween projected ealaxies which are truly within a sphere of radius RB. aud hose which are bevoud but happen to fall along the line of sight.," In physical terms, this follows because we cannot distinguish between projected galaxies which are truly within a sphere of radius $R$, and those which are beyond but happen to fall along the line of sight."
+ Considerations about the plivsical xoperties of the tracer mav help reduce he confusion. but a bias remains in the estimated mass when the sample is truncated in projected radi.," Considerations about the physical properties of the tracer may help reduce the confusion, but a bias remains in the estimated mass when the sample is truncated in projected radii."
+ Extrapolation provides sole neans of correction., Extrapolation provides some means of correction.
+ Note that extrapolation has a different mieauiug depending on what the true profile is., Note that extrapolation has a different meaning depending on what the true profile is.
+" Specifically the boundary couditious (exponential spliues. οσο spline, truncation at two or five tunes the last CASTILE radius) will make a difference in the recovered profile."," Specifically the boundary conditions (exponential splines, edge spline, truncation at two or five times the last measured radius) will make a difference in the recovered profile."
+ Since the completeness of our redshift catalogue decreases with increasing radius. we restrict our analysis to the inner region of the cluster within 1000 arcsec (1.62 Moc at the cluster redshift). where this catalogue is fairly complete (for R< 15).," Since the completeness of our redshift catalogue decreases with increasing radius, we restrict our analysis to the inner region of the cluster within 1000 arcsec (1.62 Mpc at the cluster redshift), where this catalogue is fairly complete for $\leq 18$ )."
+ Iu practice we check that all lass estimates converge to the same total mass within the error bars., In practice we check that all mass estimates converge to the same total mass within the error bars.
+ The equilibrium of an isotropic stationary spherical ealactic cluster obevs Jeaus equation: where ο) is the gravitational potential generated by all the types of matter. ie. stellar matter. X-ray ciuitting plasma auc uusecu-matter. p(r) the density of galaxies in the cluster aud σι) the radial velocity dispersion.," The equilibrium of an isotropic stationary spherical galactic cluster obeys Jeans' equation: where $\psi(r)$ is the gravitational potential generated by all the types of matter, i.e. stellar matter, X-ray emitting plasma and unseen-matter, $\rho(r)$ the density of galaxies in the cluster and $\sigma_r(r)$ the radial velocity dispersion."
+ can be applied locally to assess the cumulative dynamical lass profile., can be applied locally to assess the cumulative dynamical mass profile.
+ The surface density of galaxies is related to the density via an Abel transfor: where X(R) is the projected galaxy density and R the projected radius as measured on the sky., The surface density of galaxies is related to the density via an Abel transform: where $\Sigma(R)$ is the projected galaxy density and $R$ the projected radius as measured on the sky.
+ Similarly the line of sight velocity dispersion 0. is related. to the intrinsic radial velocity dispersion. or). via the Abel transform (or projection) Note that “(Ror is the projected kinetic energv density divided by three (corresponding to oue degree of freedom) aud ptr)o? the kiuetic energy deusity divided by three.," Similarly the line of sight velocity dispersion $\sigma_{p}^{2}$ is related to the intrinsic radial velocity dispersion, $\sigma_r^{2}(r)$, via the Abel transform (or projection) Note that $\Sigma(R) \sigma_{p}^{2} $ is the projected kinetic energy density divided by three (corresponding to one degree of freedom) and $\rho(r) \sigma_r^{2} $ the kinetic energy density divided by three."
+ Iuvertiug Eqs. (, Inverting Eqs. (
+3)-(1) iuto vields: Therefore. assuming we lave access to estimators for X aud XoD5. the euuulative mass distribution follows.,"3)-(4) into yields: Therefore, assuming we have access to estimators for $\Sigma$ and $\Sigma \sigma_{p}^{2}$, the cumulative mass distribution follows."
+ The Balicall Tremaine (1981) mass estimator for test particles around a point mass assumescompleteness and isotropy and is elven by: Iu a nutshell. given that formally the inverse of is," The Bahcall Tremaine (1981) mass estimator for test particles around a point mass assumescompleteness and isotropy and is given by: In a nutshell, given that formally the inverse of is"
+that a flat Dux spectrum in wavelength space implies a=2 when considered in frequency space).,that a flat flux spectrum in wavelength space implies $\alpha=2$ when considered in frequency space).
+ Fitting the (galaxy-subtracted) blue spectrum of the nucleus of 6-80-15. we conclude that aon. lies in the range 4.55.," Fitting the (galaxy-subtracted) blue spectrum of the nucleus of $-$ 6-30-15, we conclude that $\alpha_{\rm obs}$ lies in the range 4.5–5."
+ Assuming that he intrinsic optical spectrum of 6-30-15 is similar to hat found in unreddened Seyfert 1 nuclei. the reddening of the optical continuum source £(2V) is in the range ," Assuming that the intrinsic optical spectrum of $-$ 6-30-15 is similar to that found in unreddened Seyfert 1 nuclei, the reddening of the optical continuum source $E(B-V)$ is in the range 0.65–0.78."
+The reddening of the optical continuum. source is consistent with the lower end of the reddening derived rom the La /ll? Balmer ratio., The reddening of the optical continuum source is consistent with the lower end of the reddening derived from the $\alpha$ $\beta$ Balmer ratio.
+ Assuming that the optical Continuum is associated. with an accretion disk embedded inside the BLK. this result suggests that little dust is present )etween the accretion disk and the region where the bulk of he the broad line photons are emitted.," Assuming that the optical continuum is associated with an accretion disk embedded inside the BLR, this result suggests that little dust is present between the accretion disk and the region where the bulk of the the broad line photons are emitted."
+" Initially, we shall suppose that. the. observed line emitting region is seen directly rather than via scattered photons."," Initially, we shall suppose that the observed line emitting region is seen directly rather than via scattered photons."
+ Furthermore. suppose that the UV continuum and Dine emission are subject to the same extinction. as the optical non-stcllar continuum/line regions.," Furthermore, suppose that the UV continuum and line emission are subject to the same extinction as the optical non-stellar continuum/line regions."
+ Since we have determined the optical reddening to be in the range HRBVV)=0.61L09. we can deredden the line in order to estimate its intrinsie (1.0. dereddened) ux.," Since we have determined the optical reddening to be in the range $E(B-V)=0.61-1.09$, we can deredden the line in order to estimate its intrinsic (i.e. dereddened) flux."
+ Thus.direclly we constrain the intrinsic line [ux to lie in the range where we have included the I-0 errors in the observed flux and used the UV extinction law o£ Osterbrock (1989).," Thus, we constrain the intrinsic line flux to lie in the range where we have included the $\sigma$ errors in the observed flux and used the UV extinction law of Osterbrock (1989)."
+ This is à rather large line lux. corresponding to an isotropic luminosity of 3.LOeres or greater in theIv. line alone.," This is a rather large line flux, corresponding to an isotropic luminosity of $3\times 10^{42}\ergps$ or greater in the line alone."
+ To quantitatively assess how large this line Bux is. consider the lower end of this range corresponding to £(51322 0.61.," To quantitatively assess how large this line flux is, consider the lower end of this range corresponding to $E(B-V)=0.61$ ."
+ For this reddening. the deredcdened L3 [ux is leading to a lower limit on the intrinsic /112 flux ratio of 15.," For this reddening, the dereddened $\beta$ flux is leading to a lower limit on the intrinsic $\beta$ flux ratio of 15."
+ This ratio is very sensitive to the reddening assumed ancl cangreatly exeee js value i£. (PeV) 0.61., This ratio is very sensitive to the reddening assumed and can exceed this value if $E(B-V)>0.61$ .
+ In unredcdened: Sevfert nuclei. this ratio 1s often significantly. smaller.," In unreddened Seyfert nuclei, this ratio is often significantly smaller."
+ For cxample. in the AGN Watch Campaign on NGC 3783. the intrinsic Ll? flux ratio is 10 (Reichert et al.," For example, in the AGN Watch Campaign on NGC 3783, the intrinsic $\beta$ flux ratio is $\sim 10$ (Reichert et al."
+ 1994: Stirpe et al., 1994; Stirpe et al.
+ 1994)., 1994).
+ Similarly. the £13 Hux ratios found during the monitoring campaigns on NGC 4151. (Crenshaw et al.," Similarly, the $\beta$ flux ratios found during the monitoring campaigns on NGC 4151 (Crenshaw et al."
+ 1996: WKaspi et al., 1996; Kaspi et al.
+ 1996) and NGC 5548 (Ixorista et al., 1996) and NGC 5548 (Korista et al.
+ L995) are ~7 and ~9. respectively.," 1995) are $\sim 7$ and $\sim 9$, respectively."
+ We must conclude that the line flux is unusually high compared. with the optical line Duxes. or that one of our assumptions has broken down.," We must conclude that the line flux is unusually high compared with the optical line fluxes, or that one of our assumptions has broken down."
+ There are three possible wavs that our above argument might be Dawed., There are three possible ways that our above argument might be flawed.
+ First. source variability during the 9 months separating the UV. and optical observations may produce an apparently unusual lino ratio. even if the intrinsic line ratio is normal.," First, source variability during the 9 months separating the UV and optical observations may produce an apparently unusual line ratio, even if the intrinsic line ratio is normal."
+ In our minimum. reddening case (01=061). only mild variability (30 per cent over 9 months) is required to make the observed vI/1E2 ratio of 15 consistent. with he that seen in other objects.," In our minimum reddening case $E(B-V)=0.61$ ), only mild variability $\sim 30$ per cent over 9 months) is required to make the observed $\beta$ ratio of 15 consistent with the that seen in other objects."
+ As one postulates higher reddening values. the more extreme is the inferred intrinsic ine ratio and the more violent the variability needed.," As one postulates higher reddening values, the more extreme is the inferred intrinsic line ratio and the more violent the variability needed."
+ Secondly. the reddening towards the high-ionization BLK (including the line emitting region) may be cillerent han that towards the low-ionization DLIt (which includes he Balmer line emitting region).," Secondly, the reddening towards the high-ionization BLR (including the line emitting region) may be different than that towards the low-ionization BLR (which includes the Balmer line emitting region)."
+ Whilst this is clearly a viable possibility. (and one can imagine central-engine ecometries that produce such an clleet) there is no precedent or the high-ionization DLIt to be less reddened than the ow-ionization BLR., Whilst this is clearly a viable possibility (and one can imagine central-engine geometries that produce such an effect) there is no precedent for the high-ionization BLR to be less reddened than the low-ionization BLR.
+ Vhirdly. some fraction of the photons rom the BL might be scattered around the material responsible for the extinction.," Thirdly, some fraction of the photons from the BLR might be scattered around the material responsible for the extinction."
+ Lf the scattering fraction is wavelength independent. (c.g. electron. scattering). the scattering will tend to preferentially enhance the UV relative to the optical due to the fact that the Dux is heavily reddened.," If the scattering fraction is wavelength independent (e.g. electron scattering), the scattering will tend to preferentially enhance the UV relative to the optical due to the fact that the flux is heavily reddened."
+ Since we know scattering to be an. important process in some other Sevfert. nuclei. we now explore this last possibility in more detail.," Since we know scattering to be an important process in some other Seyfert nuclei, we now explore this last possibility in more detail."
+ Suppose that the intrinsic UVoptical line spectrum is similar to that ofNGC 3783. witha Dux ratio of 10.," Suppose that the intrinsic UV/optical line spectrum is similar to that of NGC 3783, with a $\beta$ flux ratio of 10."
+ We can write the observed fluxes of both of these lines. ην. as a sum of the direct (extinguished) Εικ and the scattered lux which is assumed. not to suller any extinction bevond that due to Galactic material.," We can write the observed fluxes of both of these lines, $F_{\rm obs}$ , as a sum of the direct (extinguished) flux and the scattered flux which is assumed not to suffer any extinction beyond that due to Galactic material."
+ I£ f is the scattering fraction. then we have where b is a parameter dependent on the extinction aw used.," If $f$ is the scattering fraction, then we have where $b$ is a parameter dependent on the extinction law used."
+ The standard. interstellar extinction. curve of Osterbrock (1989) gives 6=3.2 lor 1A1549 and 6=1.45 or 113., The standard interstellar extinction curve of Osterbrock (1989) gives $b=3.2$ for $\lambda 1549$ and $b=1.45$ for $\beta$.
+ The first term on the right hand. side of equation (8) represents the scattered. flux including the ellects of extinction by Galactic material., The first term on the right hand side of equation (8) represents the scattered flux including the effects of extinction by Galactic material.
+ We take οV)=0.06 )orriman 1989).," We take $E_{\rm
+Gal}(B-V)=0.06$ (Berriman 1989)."
+ The second term ofequation (8) gives the contribution due to the direct (extinguished) Dux., The second term of equation (8) gives the contribution due to the direct (extinguished) flux.
+ Dividing hese equations for IvVAI549 and Ll? gives à. relation tween the required scattering fraction f and the total line-ol-sight reddening £(2V7)., Dividing these equations for $\lambda 1549$ and $\beta$ gives a relation between the required scattering fraction $f$ and the total line-of-sight reddening $E(B-V)$.
+ This relationship is shown in Fig., This relationship is shown in Fig.
+ 5 for interesting values of £(2BV)., 5 for interesting values of $E(B-V)$.
+ Ht can be seen hat scattering fraction of between 15 per cent (depending on the total reddening) are required in order to make the observed line ratios consistent with an intrinsic line ratio of 10., It can be seen that scattering fraction of between 1–5 per cent (depending on the total reddening) are required in order to make the observed line ratios consistent with an intrinsic $\beta$ line ratio of 10.
+ X-rays are thought to be produced in the very central regions of the accretion disk., X-rays are thought to be produced in the very central regions of the accretion disk.
+ Phe absorption of those X-rays by line- material is readilyobservable at soft. X-ray energies ancl provides important information on the surroundings of the active nucleus., The absorption of those X-rays by line-of-sight material is readilyobservable at soft X-ray energies and provides important information on the surroundings of the active nucleus.
+ The discussion of N-ray. absorption is, The discussion of X-ray absorption is
+constraints that can be placed on these quantities in section 4.,constraints that can be placed on these quantities in section \ref{sec:res}.
+ We will assume that the hot eas is istributecl [ike the dark matter. ddensity. proportional to r7.," We will assume that the hot gas is distributed like the dark matter, density proportional to $r^{-2}$."
+ For our kinematic study we would like to consider hot gas associated with each galaxy in the halo. while in the galaxy formation model. all hot. gas is assumed to be associated with the central galaxy.," For our kinematic study we would like to consider hot gas associated with each galaxy in the halo, while in the galaxy formation model, all hot gas is assumed to be associated with the central galaxy."
+" To do this we assume that the hot gas in each subhalo is related by a parameter fous, to the total amount ofdark matter in the subhalo.", To do this we assume that the hot gas in each subhalo is related by a parameter $f_{sub}$ to the total amount of dark matter in the subhalo.
+" Thus the mass of hot eas. mp, in à given subhalo is where Won ds the dark matter mass of that subhalo 318and. AM, anc Mj, ave the total mass in hot gas and dark matter respectively. summing over all subhalos and the parent halo."," Thus the mass of hot gas, $m_{hg}$ in a given subhalo is where $m_{dm}$ is the dark matter mass of that subhalo and $M_{hg}$ and $M_{dm}$ are the total mass in hot gas and dark matter respectively, summing over all subhalos and the parent halo."
+ The remaining hot gas is hen associated with the parent halo., The remaining hot gas is then associated with the parent halo.
+ In. our realizations we find that subhalos can contain up to 2/3 of the halos mass. requiring foi. to have values between 0 and. 1.5.," In our realizations we find that subhalos can contain up to $2/3$ of the halos mass, requiring $f_{sub}$ to have values between $0$ and $1.5$."
+" In he case that fi,=0. there is no hot gas associated with he subhalos and our high-ion model is similar to the one explored in ?.."," In the case that $f_{sub} = 0$, there is no hot gas associated with the subhalos and our high-ion model is similar to the one explored in \citet{wp:00b}."
+ We note that this is far from ai model of the werodyvnamics of the hot σας in the halo. which is rather dillieult to model especially including the poorly understood ohvsies of supernova feedback.," We note that this is far from a model of the hydrodynamics of the hot gas in the halo, which is rather difficult to model especially including the poorly understood physics of supernova feedback."
+ A successful model of eas in he halo of the Milkv. Was. where there is plentiful data. still cludes researchers (?).. so we must accept that at the oresent moment we can only hope to sketch a rudimentary ramework of the situation at high redshilt.," A successful model of gas in the halo of the Milky Way, where there is plentiful data, still eludes researchers \citep{cbr:02}, so we must accept that at the present moment we can only hope to sketch a rudimentary framework of the situation at high redshift."
+ For the purpose of simulating spectra we assume tha a fraction feri of the hot gas is no longer at the viria emperature but has cooled into pressure confined. clouds and is. photoionizecl.. with. a temperature =5rOF1 deglx.- We use 15 clouds to properly sample the velocity field. zux assume clouds have random velocities. cv. proportiona to the halo's circular velocity (see Table 1)).," For the purpose of simulating spectra we assume that a fraction $\fciv$ of the hot gas is no longer at the virial temperature but has cooled into pressure confined clouds and is photoionized with a temperature $\approx 5 \times 10^4 \deg$ K. We use 15 clouds to properly sample the velocity field and assume clouds have random velocities, $\sigma_{CIV}$, proportional to the halo's circular velocity (see Table \ref{tab:mod}) )."
+ We compare our mocel only to the CIV data which bes samples the hottest phase of the gas., We compare our model only to the CIV data which best samples the hottest phase of the gas.
+ The fact that there are some cdillerences in the high-ion kinematics between ions (7.e.g.SLIVvs.€IV) implies variations in the ionization state or metallicity of the σας.," The fact that there are some differences in the high-ion kinematics between ions \citep[e.g.\, Si\,IV vs.\, C\,
+IV]{pw:02} implies variations in the ionization state or metallicity of the gas."
+ Future modelling of this may help us to understand metallicity. density. ancl ionization eracdients in the hot gas.," Future modelling of this may help us to understand metallicity, density and ionization gradients in the hot gas."
+ AMini-halos (ej;35kms *) are usually not. considered in discussions of galaxy formation because their expected luminosities make them cdillieult to observe., Mini-halos $\vvir \le 35 \kms$ ) are usually not considered in discussions of galaxy formation because their expected luminosities make them difficult to observe.
+ However. ? have shown that there may be sullicient cold eas in mini-halos at high redshift to comprise the LL systems at those recdshifts.," However, \citet{am:98} have shown that there may be sufficient cold gas in mini-halos at high redshift to comprise the LL systems at those redshifts."
+ Reeenthy the issue of mini-halos has also been widely aclcressed cause Of what has been named theproblem: the fact that N-bodsy simulations expect hundreds of mini-subhalos in a Alilky Way sized halo (?2).. but only a couple dozen or fewer are observed.," Recently the issue of mini-halos has also been widely addressed because of what has been named the; the fact that N-body simulations expect hundreds of mini-subhalos in a Milky Way sized halo \citep{klypin:99b,moore:99a}, but only a couple dozen or fewer are observed."
+ While many exotic solutions to this problem have been proposed. the simplest is that the extragalactic UV background suppresses the accretion anc cooling of eas in these mini-halos. and therefore stars form in only the small number that managed to collapse before the epic of reionization.," While many exotic solutions to this problem have been proposed, the simplest is that the extragalactic UV background suppresses the accretion and cooling of gas in these mini-halos, and therefore stars form in only the small number that managed to collapse before the epic of reionization."
+ Models of this scenario (2???) have been fairly successful in explaining why only a small fraction of mini-subhalos would have formed stars.," Models of this scenario \citep{bkw:00,benson:02,somer:02} have been fairly successful in explaining why only a small fraction of mini-subhalos would have formed stars."
+ ‘To explore the contribution of mini-halos to LL svstems we use thesquefehing model of ?. to determine the amount of hot and cold gas in mini and low-mass (35kmsves>5Okms +) halos., To explore the contribution of mini-halos to LL systems we use the model of \citet{somer:02} to determine the amount of hot and cold gas in mini and low-mass $35 \kms > \vvir > 50 \kms$ ) halos.
+ This model uses a fitting function found by ? in hyclrodyvnamical simulations to determine the amount of gas that can accrete onto a dark matter halo in the presence of an ionizing field., This model uses a fitting function found by \citet{gned:00} in hydrodynamical simulations to determine the amount of gas that can accrete onto a dark matter halo in the presence of an ionizing field.
+ We assume that gas will only contract to density profiles steeper than the dissipationless dark matter., We assume that gas will only contract to density profiles steeper than the dissipationless dark matter.
+ Therefore to explore themaximal contribution of these halos to LL systems. we take their projected. density. profiles to go as Rt like the cold gas disks and hot gas in galaxies described above (equation 1)).," Therefore to explore the contribution of these halos to LL systems, we take their projected density profiles to go as $R^{-1}$ like the cold gas disks and hot gas in galaxies described above (equation \ref{eq:mestel}) )."
+ We truncate the gas at the virial radius or we consider cases where the gas is truncated at a specified column density. δεν," We truncate the gas at the virial radius or we consider cases where the gas is truncated at a specified column density, $N_t$."
+ In the latter case the truncation radius can then be evaluated by where mg is the mass of the hydrogen atom., In the latter case the truncation radius can then be evaluated by R_t = where $m_H$ is the mass of the hydrogen atom.
+ We pass multiple random lines of sight through each halo to produce Monte-Carlo realizations of spectra in the manner described in paper EL. Enough realizations are performed to eive us at least 10.000 sight lines with DLA systems in them.," We pass multiple random lines of sight through each halo to produce Monte-Carlo realizations of spectra in the manner described in paper I. Enough realizations are performed to give us at least 10,000 sight lines with DLA systems in them."
+ Figure 2 shows the distribution ofCIV column densities, Figure \ref{fig:civ} shows the distribution of CIV column densities
+was attained in the disk in at least two locations. at the outer shock frout and just behind the peak of the specific augular momentum distribution.,"was attained in the disk in at least two locations, at the outer shock front and just behind the peak of the specific angular momentum distribution."
+ Those regions were subject to the rotationa instability., Those regions were subject to the rotational instability.
+ dHdi<0 caused the average angular nomentun «7- of the disk to increase. auk rence the period aud the amplitude of the shock oscillation to increase too.," $dl/dr<0$ caused the average angular momentum $<l>$ of the disk to increase, and hence the period and the amplitude of the shock oscillation to increase too."
+ This is less perceptible or lower a aud the shock oscillation achieve quasi-saturatiou. but for a=0.1 the shock wet outside the computation domain.," This is less perceptible for lower $\alpha$ and the shock oscillation achieved quasi-saturation, but for $\alpha=0.1$ the shock went outside the computation domain."
+ We repeater he simulation with a=0.3 (not presented iu he paper) and iu this case too the shock went outside the domain. although formation of ier souic point and oscillation of the two shocks were observed too.," We repeated the simulation with $\alpha=0.3$ (not presented in the paper) and in this case too the shock went outside the domain, although formation of inner sonic point and oscillation of the two shocks were observed too."
+ Iu case of multidineusional sinulatious. a part of the post shock matter would have ejected along the vertical direction in the form: of winds. which would have carried away a part of the aneular momentum. such that the increase of |oc inay have been arrested for higher o.," In case of multi-dimensional simulations, a part of the post shock matter would have ejected along the vertical direction in the form of winds, which would have carried away a part of the angular momentum, such that the increase of $<l>$ may have been arrested for higher $\alpha$."
+ This would have meant that the shock oscillation may saturate for a2O14., This would have meant that the shock oscillation may saturate for $\alpha \geq 0.1$.
+ Hence. we conjecture that the non-saturatiou of shock oscillation for a<O0.1. could be an artifact of one-dimensional siauulation.," Hence, we conjecture that the non-saturation of shock oscillation for $\alpha\gsim 0.1$, could be an artifact of one-dimensional simulation."
+ We will test it in a future work using multi-dimensional simulations., We will test it in a future work using multi-dimensional simulations.
+ SL was supported in part bv Basic Science Research Program through the National Research Foundation of Norea (NRE) fuuded by the Ministry of Education. Science and Techuology (2010-0001738).," \acknowledgments{SL was supported in part by Basic Science Research Program through the National Research Foundation of Korea (NRF) funded by the Ministry of Education, Science and Technology (2010-0004738)."
+ DR was supported in part by National Research Foundation of Ίοτοα. through erat IRE-2007-311-C'O0020., DR was supported in part by National Research Foundation of Korea through grant KRF-2007-341-C00020.
+White-light flares (WLFs) are rare energetic events characterized bv a visible continuum enhancement (o a few or tens of percent. which imposes strict constraints on the modeling of solar [lares in terms of energy release and (ransport processes in (he impulsive phase.,"White-light flares (WLFs) are rare energetic events characterized by a visible continuum enhancement to a few or tens of percent, which imposes strict constraints on the modeling of solar flares in terms of energy release and transport processes in the impulsive phase."
+ According to the spectral features. two tvpes of WLFs have been proposed," According to the spectral features, two types of WLFs have been proposed"
+"In Fig. 4,,","In Fig. \ref{fig:SED},"
+ we show the effect of dust extinction in the spectral energy distribution (SED) of GRBs., we show the effect of dust extinction in the spectral energy distribution (SED) of GRBs.
+" The effect is significant in the G band (~5x1014 Hz), which will considerably decrease the detection rate of optical GRBs, mainly at high-z. For high-z GRBs, much of the optical and near-infrared light will be absorbed by the Lya forest, which provides a powerful tool to probe the reionization era (Miralda-2011)."," The effect is significant in the G band $\sim 5 \times 10^{14}$ Hz), which will considerably decrease the detection rate of optical GRBs, mainly at high-z. For high-z GRBs, much of the optical and near-infrared light will be absorbed by the $Ly\alpha$ forest, which provides a powerful tool to probe the reionization era \citep{Miralda1998,Mesinger2008,Ciardi2011}."
+ The GRB 050904 at z — 6.3 was the first attempt to probe the IGM through GRBs at the epoch of reionization by using the damping wing at wavelengths redward of the Lyman break (Totanietal.2006;Kawaietal.2006).," The GRB 050904 at z = 6.3 was the first attempt to probe the IGM through GRBs at the epoch of reionization by using the damping wing at wavelengths redward of the Lyman break \citep{Totani2006,Kawai2006}."
+". The absorption by the neutral IGM can be approximated by (McQuinnetal.2008) where vq is the rest frame of the Lya line, Ry represents the size of an HII region surrounded by an IGM with neutral fraction ry], Zpost is the redshiftof the GRB host galaxy, and H(z) is the Hubble parameter for a ACDM cosmology."," The absorption by the neutral IGM can be approximated by \citep{McQuinn2008} where $\nu_{\alpha}$ is the rest frame of the $\rm Ly{\alpha}$ line, $R_b$ represents the size of an HII region surrounded by an IGM with neutral fraction $x_{\rm HI}$, $z_{host}$ is the redshiftof the GRB host galaxy, and H(z) is the Hubble parameter for a $\Lambda$ CDM cosmology."
+" To estimate xyz, we use the prescription detailed in deSouzaetal.(2011) (see Fig.1)."," To estimate $x_{\rm HI}$ , we use the prescription detailed in \citet{rafael2011} (see Fig.1)."
+" The optical depth of the damped Lya absorber (DLA), TDLA=Νηισα[νου»(1+Znost)], can be computed by (e.g., Barkana&Loeb 2004)) where voy,=6/λοὺς, (1+Zobs)=Aobs/Aa, Nur is the total column density of HI in the host galaxy, and λους=AA+λα(12)."," The optical depth of the damped $\alpha$ absorber (DLA), $\tau_{DLA}=N_{\rm HI} \sigma_{\alpha}[\nu_{obs}(1+z_{host})]$, can be computed by (e.g., \citealt{Barkana2004}) ) where $\nu_{obs} = c/\lambda_{obs}$, $(1+z_{obs}) = \lambda_{obs}/\lambda_{\alpha}$, $N_{\rm HI}$ is the total column density of $\rm HI$ in the host galaxy, and $\lambda_{obs}= \Delta\lambda+\lambda_{\alpha}(1+z)$."
+" Nur is randomly chosen assuming a cumulative distribution function scaling as Ng, between 1015 and 10?!5em-7? (see Chenetal.2007;McQuinn 2008))."," $N_{\rm HI}$ is randomly chosen assuming a cumulative distribution function scaling as $N_{\rm HI}^{0.3}$ , between $10^{18}$ and $10^{21.5} cm^{-2}$ (see \citealt{Chen2007,McQuinn2008}) )."
+" For each event, Ry is chosen from a lognormal distribution between 1-100 Mpc motivated by a visual inspection in Fig."," For each event, $R_b$ is chosen from a lognormal distribution between 1-100 Mpc motivated by a visual inspection in Fig."
+ 5 from McQuinnetal.(2008)., 5 from \citet{McQuinn2008}.
+. The mock sample is generated by the Monte Carlo method assuming different probability distribution functions (PDF) for each quantity as explained below., The mock sample is generated by the Monte Carlo method assuming different probability distribution functions (PDF) for each quantity as explained below.
+ The medium density n is randomly chosen from a flat distribution within 0.1— lem~?., The medium density $n$ is randomly chosen from a flat distribution within $0.1-1 cm^{-3}$ .
+ We generatethe GRB events randomly in redshift with a PDF given by Eq. (2))., We generatethe GRB events randomly in redshift with a PDF given by Eq. \ref{dngrbtrue}) ).
+ The probability of a given GRB appearingat redshift z is, The probability of a given GRB appearingat redshift $z$ is
+wavelength in the case of broad band imaging).,wavelength in the case of broad band imaging).
+" See, for example, Szeifert(2002)."," See, for example, \cite{szeifert}."
+". Finally, we have investigated the effects produced by a deviation AG from the nominal phase retardance of an HWP (7)."," Finally, we have investigated the effects produced by a deviation $\Delta\beta$ from the nominal phase retardance of an HWP $\pi$ )."
+" Using the general expression of the Mueller matrix for a retarder forexampleGoldstein2003),, the normalized flux ratios(see turn out to be from which it is clear that the measured linear polarization depends on the circular polarization of the input signal."," Using the general expression of the Mueller matrix for a retarder \citep[see for example][]{goldstein}, the normalized flux ratios turn out to be from which it is clear that the measured linear polarization depends on the circular polarization of the input signal."
+" For V=0 and AG <10° the corresponding absolute error of the polarization degree is less than0.05%,, while the outcome on the polarization angle is negligible."," For $V$ =0 and $\Delta\beta \leq$ $^\circ$ the corresponding absolute error of the polarization degree is less than, while the outcome on the polarization angle is negligible."
+ For V 40 the exact effect depends on the ratio between circular and linear polarization degrees., For $V\neq$ 0 the exact effect depends on the ratio between circular and linear polarization degrees.
+" For example, for A =10° and Q=UV =0.01, the absolute error of the computed polarization degree is about for N=4, which decreases to for N=16."," For example, for $\Delta=$ $^\circ$ and $Q=U=V=$ 0.01, the absolute error of the computed polarization degree is about for $N$ =4, which decreases to for $N$ =16."
+ It is worthwhile noting that this defect would be detected by a Fourier as a component with a period 7 and whose intensity is analysis|VsinA]., It is worthwhile noting that this defect would be detected by a Fourier analysis as a component with a period $\pi$ and whose intensity is $|\bar{V}\sin\Delta|$.
+" Typically the analyzer is followed by additional optics, like filters, grisms and camera lenses which, due to their possible tilt with respect to the optical axis, may behave as poor linear polarizers."," Typically the analyzer is followed by additional optics, like filters, grisms and camera lenses which, due to their possible tilt with respect to the optical axis, may behave as poor linear polarizers."
+" In the most probable case where the polarization is produced by transmission (see also next section), the properties of post-analyzer (PA) components can be described by the following approximate Mueller matrix: where A zl, C zl, c=cos2y, s=sin2y and 9 is the polarization angle can change across the field of while B is (whichrelated to the polarization degree introducedview) by the PA optics."," In the most probable case where the polarization is produced by transmission (see also next section), the properties of post-analyzer (PA) components can be described by the following approximate Mueller matrix: where $A\approx$ 1, $C\approx$ 1, $c=\cos 2\varphi$, $s=\sin 2\varphi$ and $\varphi$ is the polarization angle (which can change across the field of view) while $B$ is related to the polarization degree introduced by the PA optics."
+" This expression can be deduced from the general formulation (seeKeller2002,Eq. after applying the usual matrix rotation (see4.63)forexampleKeller2002,Eq.2.5).", This expression can be deduced from the general formulation \citep[see][Eq. 4.63]{keller} after applying the usual matrix rotation \citep[see for example][Eq. 2.5]{keller}.
+". If Sy=Qo,Uo, is the input Stokes vector, the effect of PA optics (Jo,Vo)can be evaluated computing the Stokes vectors that correspond to the ordinary and extraordinary beams produced by the WP, transforming them using the operator Mp, and using the resulting intensity components to compute the normalized flux differences F;."," If $\vec{S_0}=(I_0,Q_0,U_0,V_0)$ is the input Stokes vector, the effect of PA optics can be evaluated computing the Stokes vectors that correspond to the ordinary and extraordinary beams produced by the WP, transforming them using the operator $M_{PA}$ and using the resulting intensity components to compute the normalized flux differences $F_i$."
+" After simple calculations, one arrives at the following expression: were we have assumed that |B|<1, i.e. that the linear polarization induced by the PA optics is small."," After simple calculations, one arrives at the following expression: were we have assumed that $|B|\ll 1$, i.e. that the linear polarization induced by the PA optics is small."
+" Given this expression, it is clear that the redundancy in the HWP positions (V=4, 8, 16) eliminates this problem, since the additive term DBcos2« is not modulated by the HWP rotation, while for N—2 the derived polarization degree and angle would be affected, possibly in a severe way."," Given this expression, it is clear that the redundancy in the HWP positions $N$ =4, 8, 16) eliminates this problem, since the additive term $B\cos 2\varphi$ is not modulated by the HWP rotation, while for $N$ =2 the derived polarization degree and angle would be affected, possibly in a severe way."
+" If the optimal HWP angle set has been used, it is easy to verify that Bcos2p=F,/N, which is identical to Eq. 16.."," If the optimal HWP angle set has been used, it is easy to verify that $B\cos 2\varphi = \sum_{i=0}^{N-1} F_i/N$, which is identical to Eq. \ref{eq:a0}. ."
+" This means that,aire in a first approximation, it is not possible to disentangle between an imperfect WP and the presence of polarization in the PA optics."," This means that, in a first approximation, it is not possible to disentangle between an imperfect WP and the presence of polarization in the PA optics."
+" Therefore, the fact that αρ~0.4% in Fig. 3,,"," Therefore, the fact that $a_0\simeq$ in Fig. \ref{fig:fourier},"
+ can actually be attributed to both kinds of problems., can actually be attributed to both kinds of problems.
+" The PA optics effect becomes definitely stronger when these include highly tilted components, like the grisms."," The PA optics effect becomes definitely stronger when these include highly tilted components, like the grisms."
+" This is very well illustrated by the two examples of Fig. 4,,"," This is very well illustrated by the two examples of Fig. \ref{fig:postanal},"
+" where we show the results obtained using VLT-FORS1 archival data of a highly polarized star (Vela 1 95, a=09:06:00, --4τ:19:00) and an unpolarized star (WD 1615-154, ὃ--a=16:17:55, 6=—15:35:51)?,, which were observed on the optical axis, where the instrumental polarization is known to be null (Szeifert2002)."," where we show the results obtained using VLT-FORS1 archival data of a highly polarized star (Vela 1 95, $\alpha$ =09:06:00, $\delta$ $-$ 47:19:00) and an unpolarized star (WD 1615-154, $\alpha$ =16:17:55, $\delta$ $-$, which were observed on the optical axis, where the instrumental polarization is known to be null \citep{szeifert}."
+. In both cases the polarization degree deduced using N=2 (central is markedly different from that derived with N=4 panels)(upper and the deviation is particularly severe for the panels)unpolarized object., In both cases the polarization degree deduced using $N$ =2 (central panels) is markedly different from that derived with $N$ =4 (upper panels) and the deviation is particularly severe for the unpolarized object.
+" As one can finally notice, the resulting values of Bcos2y show a strong wavelength dependency and are higher than at about 800 nm."," As one can finally notice, the resulting values of $B\cos
+2\varphi$ show a strong wavelength dependency and are higher than at about 800 nm."
+" It is interesting to note that Bcos2y ~0 at about 450 nm, i.e. at the wavelength where the anti-reflection coatings are optimized (see next "," It is interesting to note that $B\cos 2\varphi\sim$ 0 at about 450 nm, i.e. at the wavelength where the anti-reflection coatings are optimized (see next section)."
+"This fact, together with the marked wavelength section).dependency and the much lower level seen in broad band imaging (see Fig. 3)),"," This fact, together with the marked wavelength dependency and the much lower level seen in broad band imaging (see Fig. \ref{fig:fourier}) ),"
+ strongly suggest that the effect seen in Fig., strongly suggest that the effect seen in Fig.
+ 4 isindeed produced by the tilted surfaces of the grism., \ref{fig:postanal} isindeed produced by the tilted surfaces of the grism.
+the right side of the array.,the right side of the array.
+ This reduced the signal-to-noise ratio in the short-wavelength half of each order., This reduced the signal-to-noise ratio in the short-wavelength half of each order.
+" All data reduction procedures were carried out using ΗΑΕ,", All data reduction procedures were carried out using IRAF.
+ Subsequent images were subtracted from each other. removing the background. ancl (ος emission lines to first order.," Subsequent images were subtracted from each other, removing the background and telluric emission lines to first order."
+ A flat field correction was applied to each clilferenced image., A flat field correction was applied to each differenced image.
+ To minimize the impact of imperfectly cancelled telluric emission lines on the extracted spectrum. we rectified the images so (hat (he image of the slit ran vertically along pixel columns.," To minimize the impact of imperfectly cancelled telluric emission lines on the extracted spectrum, we rectified the images so that the image of the slit ran vertically along pixel columns."
+ The required geometric transformation was obtained [rom the original pattern of telluric emission lines using the IRAF tasks FITCOORDS and TRANSFORAL., The required geometric transformation was obtained from the original pattern of telluric emission lines using the IRAF tasks FITCOORDS and TRANSFORM.
+ Individual diflerenced images were average combined and one-dimensional spectra were extracted lor each nod position and order., Individual differenced images were average combined and one-dimensional spectra were extracted for each nod position and order.
+ Decause the spectra had been rectified. large background apertures (located on both sides of the source aperture) could be used in extracting the spectrum. producing a higher specirum.," Because the spectra had been rectified, large background apertures (located on both sides of the source aperture) could be used in extracting the spectrum, producing a higher signal-to-noise spectrum."
+ Telluric absorption features were removed by dividing the program spectrum by that ol a tellurie standard (the hot star WR 1605; spectral (ype AS). whieh was obtained and reduced in (he same wav as the V836 Tau spectrum.," Telluric absorption features were removed by dividing the program spectrum by that of a telluric standard (the hot star HR 1605; spectral type A8), which was obtained and reduced in the same way as the V836 Tau spectrum."
+ The spectrum of HIA. 1605 was mostly featureless except. [or absorption in the 75 Ply line., The spectrum of HR 1605 was mostly featureless except for absorption in the 7–5 $\beta$ line.
+ Consequently the divided spectrum has artificially enhanced hydrogen emission in this region.," Consequently, the divided spectrum has artificially enhanced hydrogen emission in this region."
+ To correct [or the dillerence in airmass between the object and the standard spectra. we used the RAF task TELLURIC to perform the division.," To correct for the difference in airmass between the object and the standard spectra, we used the IRAF task TELLURIC to perform the division."
+ Telluric absorption lines were used for the wavelength calibration., Telluric absorption lines were used for the wavelength calibration.
+" An approximate [τις calibration was obtained by acquiring additional spectra of both. V836 Tau and HIX1605 with the 0.72"" NIRSPEC slit.", An approximate flux calibration was obtained by acquiring additional spectra of both V836 Tau and HR1605 with the $^{\prime\prime}$ NIRSPEC slit.
+ Given the poor seeing. slit losses are likely to have been signilicant.," Given the poor seeing, slit losses are likely to have been significant."
+ The short integrations used [or the V336 Tau observations result in a formal 1-0 error of in the flux calibration of each order., The short integrations used for the V836 Tau observations result in a formal $\sigma$ error of in the flux calibration of each order.
+ The possibility of slit losses contributes additional uncertainty in the fIux., The possibility of slit losses contributes additional uncertainty in the flux.
+ Nevertheless. the resulting continuum fluxes for each order agree to ~15% with the 4.5jm Spitzer/IRAC flux reported for V836 Tau (Padgett et 22006: [Fy=0.127 Jy).," Nevertheless, the resulting continuum fluxes for each order agree to $\sim 15$ with the $4.5\micron$ Spitzer/IRAC flux reported for V836 Tau (Padgett et 2006; $F_{4.5}=0.127$ Jy])."
+ For comparison. our calibrated continuum flux of 0.128. Jv. the average of the continuum level in orders 15 and 16. is ~80% brighter than the value of ~0.07 Jv that we reported previously for V836 Tau in the same wavelength region based on NIRSPEC data taken in 2001 (Najita et 22003).," For comparison, our calibrated continuum flux of 0.128 Jy, the average of the continuum level in orders 15 and 16, is $\sim 80$ brighter than the value of $\sim 0.07$ Jy that we reported previously for V836 Tau in the same wavelength region based on NIRSPEC data taken in 2001 (Najita et 2003)."
+ It is also fainter than the 0.18 Jy continuum Παν at d.78/0n (hat is implied by extrapolating the Spitzer/IRS spectrum of νου Tau (Furlan et 22006) to that wavelength., It is also fainter than the 0.18 Jy continuum flux at $4.78\micron$ that is implied by extrapolating the Spitzer/IRS spectrum of V836 Tau (Furlan et 2006) to that wavelength.
+ Because of the relatively large uncertaintv in the flux calibration lor each order. we are unable (o determine if (here is a shallow continuum slope between the two orders.," Because of the relatively large uncertainty in the flux calibration for each order, we are unable to determine if there is a shallow continuum slope between the two orders."
+ We therefore adopted (he continuum flux slope measured by Furlan et, We therefore adopted the continuum flux slope measured by Furlan et
+Thus for the NEW profile with a limiting value of f(c)=4/3 for small e. whereas for the Moore profile. This expression diverges logarithmically as (riii σοος o zero. rellecting the fact that the Dux from a pure ke.17 cusp is infinite.,"Thus for the NFW profile with a limiting value of $f(c) = 4/3$ for small $c$, whereas for the Moore profile, This expression diverges logarithmically as ${x_{\rm min}}$ goes to zero, reflecting the fact that the flux from a pure $r^{-1.5}$ cusp is infinite."
+ In. practice. annihilation would. reduce the densitv at the centre of the halo to a finite value even or pure ΝΤ. while in real halos barvons will also alfect alo structure on small scales.," In practice, annihilation would reduce the density at the centre of the halo to a finite value even for pure CDM, while in real halos baryons will also affect halo structure on small scales."
+" Furthermore. the simulations hat indicate central cusps in halos can only probe down ο rfr,c0.01. or rire&OL (Power et 22002): odow this scale. it is possible that the slope of the profile continues to change (Lavlor Navarro 2001)."," Furthermore, the simulations that indicate central cusps in halos can only probe down to $r/r_v \simeq 0.01$, or $r/r_s \simeq 0.1$ (Power et 2002); below this scale, it is possible that the slope of the profile continues to change (Taylor Navarro 2001)."
+ Thus in what follows we will consider. Moore profiles truncated. at various inner cutolf radii. either suus=LOI (profile MI yereafter). igi=107 (profile M2 hereafter). rin=101 (profile N33 hereafter). or rnin=1/3 (profile M4 hereafter). with cores of constant density within the ο{ο racius. (," Thus in what follows we will consider Moore profiles truncated at various inner cutoff radii, either $x_{\rm min} = 10^{-10}$ (profile M1 hereafter), $x_{\rm min} = 10^{-2}$ (profile M2 hereafter), $x_{\rm min} = 10^{-1}$ (profile M3 hereafter), or $x_{\rm min} = 1/3$ (profile M4 hereafter), with cores of constant density within the cutoff radius. ("
+Adding a core mocdifies equations (14)) and (16)) slightly. out. the dillerence is negligible for .r Juin.),"Adding a core modifies equations \ref{eq:nfw_fc}) ) and \ref{eq:moore_fc}) ) slightly, but the difference is negligible for $x \gg x_{\rm min}$ .)"
+ Lhe first eutoll corresponds the scale on which annihilation alone might truncate a pure power-law density profile (Cálleánneo-toldan Moore 2000) and the third cutolf corresponds to he largest core consistent with the results of current. high-resolution simulations. while the fourth cutolf corresponds to he limit of the region within which barvons might plausibly rave Lattened out the core in the halos of massive galaxies.," The first cutoff corresponds the scale on which annihilation alone might truncate a pure power-law density profile (C\'{a}llc\'{a}nneo-Roldan Moore 2000) and the third cutoff corresponds to the largest core consistent with the results of current high-resolution simulations, while the fourth cutoff corresponds to the limit of the region within which baryons might plausibly have flattened out the core in the halos of massive galaxies."
+ As mentioned above. the simulations that provide evidence or a universal density profile for CDM. halos can only reliably resolve two decades of raclius below the virial radius.," As mentioned above, the simulations that provide evidence for a universal density profile for CDM halos can only reliably resolve two decades of radius below the virial radius."
+ While the logarithmic slope appears to converge to a value rchween 195r and. Lin the inner regions. it is clearly not a simple power-law in the outer regions. where it. llattens substantially over a decade in radius.," While the logarithmic slope appears to converge to a value between $-1.5$ and $-1$ in the inner regions, it is clearly not a simple power-law in the outer regions, where it flattens substantially over a decade in radius."
+ Thus. there remains he possibility that the inner slope of CDM halos. were it resolved to smaller radii. would continue to change in value.," Thus, there remains the possibility that the inner slope of CDM halos, were it resolved to smaller radii, would continue to change in value."
+ A shallower inner slope would help to explain the rotation curves of dark-matter dominated cwarf galaxies BBlais-Ouellette. Amram. Carignan 2001: de Blok Bosma 2002). as well as the mass clistribution of the Milkv Way within the solar circle (Binney Evans 2001).," A shallower inner slope would help to explain the rotation curves of dark-matter dominated dwarf galaxies Blais-Ouellette, Amram, Carignan 2001; de Blok Bosma 2002), as well as the mass distribution of the Milky Way within the solar circle (Binney Evans 2001)."
+ One possible form for such a profile was proposed. by ‘Taylor and Navarro (2001. TN hereafter).," One possible form for such a profile was proposed by Taylor and Navarro (2001, TN hereafter)."
+ In the outer regions it resembles the NEW or. Moore. profiles. while in the inner regions (below about log(r/r;)=1.5). its logarithmic slope decreases slowly to 0.75.," In the outer regions it resembles the NFW or Moore profiles, while in the inner regions (below about $\log(r/r_v) = -1.5$ ), its logarithmic slope decreases slowly to $-0.75$."
+ An analytic fit to the inner slope Is: where αν=rírpw.o with rpw=ονυν.," An analytic fit to the inner slope is: where $x = r/r_{\rm TN}$, with $r_{\rm TN} = 5/3\, r_{s,{\rm NFW}}$."
+ We will calculate. f(c). which is convergent and well-behavecl for this profile. from a direct. numerical integration. although it can be computed to reasonable accuracy using the fitting formula above.," We will calculate $f(c)$, which is convergent and well-behaved for this profile, from a direct numerical integration, although it can be computed to reasonable accuracy using the fitting formula above."
+ From equation (4)). the limiting value for small ο is f(c)=9/8.," From equation \ref{powlaw}) ), the limiting value for small $c$ is $f(c) = 9/8$."
+ Note that in its outer regions. the equilibrium density. profile derived by TN goes to zero at a finite radius.," Note that in its outer regions, the equilibrium density profile derived by TN goes to zero at a finite radius."
+ This assumes an equilibrium that will never be achieved in cosmological setting. however. where halos will continue to acerete material onto their outer parts.," This assumes an equilibrium that will never be achieved in cosmological setting, however, where halos will continue to accrete material onto their outer parts."
+ Thus. we assume that bevond τν=1. the profile drops olf as ke. as in the NEW and. Moore fits.," Thus, we assume that beyond $x_{\rm TN} = 1$, the profile drops off as $r^{-3}$, as in the NFW and Moore fits."
+ ligure Lo shows the densitv. profiles (top panel) ancl f(r) (bottom panel) for the six profiles discussed above. including [our versions of the Moore profile with cdilferent. constant-density cores (within radii indicated by the horizontal dashed lines).," Figure \ref{fig:1} shows the density profiles (top panel) and $f(x)$ (bottom panel) for the six profiles discussed above, including four versions of the Moore profile with different constant-density cores (within radii indicated by the horizontal dashed lines)."
+" In order to compare the dilferent profiles. they have been plotted in terms ofr,3rfr. the radius relative to the point r, at which the circular velocity. peaks. and normalised so that the density is the same at ry."," In order to compare the different profiles, they have been plotted in terms of $x_p \equiv r/r_p$, the radius relative to the point $r_p$ at which the circular velocity peaks, and normalised so that the density is the same at $r_p$."
+" In the top panel. we see that all the density. profiles aside from M4 are in excellent agreement for ur,=0.1. the region probed by current simulations."," In the top panel, we see that all the density profiles aside from M4 are in excellent agreement for $x_p \ge 0.1$, the region probed by current simulations."
+ Existing simulations are roughly compatible with any of thefits shown here., Existing simulations are roughly compatible with any of thefits shown here.
+ The dillerences in the inferred. central behaviour of halos are based on dillerent interpretations of the density. profile very close to the current numerical resolution limit., The differences in the inferred central behaviour of halos are based on different interpretations of the density profile very close to the current numerical resolution limit.
+ The bottom panel shows the effect of this uncertainty on fle) where e is measured relative to ry. as in the top panel.," The bottom panel shows the effect of this uncertainty on $f(c)$, where $c$ is measured relative to $r_p$, as in the top panel."
+ We see that in cach case. the Dux multiplier increases with concentration. rellecting the contribution [rom the dense central core.," We see that in each case, the flux multiplier increases with concentration, reflecting the contribution from the dense central core."
+ To first approximation. the dependence on ¢ is independent of profile. while the inner slope of the halo profile determines the overall normalisation of f(e).," To first approximation, the dependence on $c$ is independent of profile, while the inner slope of the halo profile determines the overall normalisation of $f(c)$."
+ For the Moore profile it is 225 times larger than for an NEW profile. while for the shallower TN. profile it is comparable to a Aloore profile with a core between and of ry in size.," For the Moore profile it is 2–25 times larger than for an NFW profile, while for the shallower TN profile it is comparable to a Moore profile with a core between and of $r_s$ in size."
+ The solid vertical lines on the plot indicate typical values of e for present-day galaxies and clusters., The solid vertical lines on the plot indicate typical values of $c$ for present-day galaxies and clusters.
+ For galaxy halos. fte) is typically 50 for an NEW profile. 200 for the TN prolile or profile M2. and. 1200 for the most extreme Moore profile. ALL.," For galaxy halos, $f(c)$ is typically 50 for an NFW profile, 200 for the TN profile or profile M2, and 1200 for the most extreme Moore profile, M1."
+ For all profiles it scales with concentration as roughly C7 Dol.in this. range of c. as indicated. by the solid. line. in. the corner of the figure.," For all profiles it scales with concentration as roughly $c^{2.4}$ in this range of $c$ , as indicated by the solid line in the corner of the figure."
+ This figure indicates the importance of the inner region, This figure indicates the importance of the inner region
+"radiation hydrodynamics (RHD) simulations, both values are determined by averaging over 10 grid cells in vertical direction, centered on the midplane.","radiation hydrodynamics (RHD) simulations, both values are determined by averaging over 10 grid cells in vertical direction, centered on the midplane."
+" In radial direction, all grid cells, which are filled with dust are taken into account when summing up the optical depth and calculating the median velocity, respectively."," In radial direction, all grid cells, which are filled with dust are taken into account when summing up the optical depth and calculating the median velocity, respectively."
+ Plotted in yellow is the result of the analytical estimate., Plotted in yellow is the result of the analytical estimate.
+ The black triangles and blue stars refer to our RHD simulations listed in Table 2.., The black triangles and blue stars refer to our RHD simulations listed in Table \ref{tab:params}.
+" In general, the simulations are in reasonable agreement with the simple analytical estimate."," In general, the simulations are in reasonable agreement with the simple analytical estimate."
+" As expected, the agreement is better for the high optical depth case."," As expected, the agreement is better for the high optical depth case."
+" For decreasing optical depth values, smaller and smaller areas of the clouds fulfill the assumption that all lines of sight through the clouds are optically thick."," For decreasing optical depth values, smaller and smaller areas of the clouds fulfill the assumption that all lines of sight through the clouds are optically thick."
+" The deviations arise from the fact that in the real hydrodynamical simulations, the clouds get disrupted by the action of the radiation pressure and gravitational forces together with the onset of instabilities, leading to overall larger sizes and, therefore, smaller column densities and higher outflow velocities."," The deviations arise from the fact that in the real hydrodynamical simulations, the clouds get disrupted by the action of the radiation pressure and gravitational forces together with the onset of instabilities, leading to overall larger sizes and, therefore, smaller column densities and higher outflow velocities."
+ In the following we summarise the main results from several parameter studies:, In the following we summarise the main results from several parameter studies:
+2004).,.
+". As we cannot accurately determine the amount of missed flux, we adopt the point source aperture correction of 0.1 and note this uncertainty."," As we cannot accurately determine the amount of missed flux, we adopt the point source aperture correction of 0.1 magnitudes, and note this systematic uncertainty."
+" In order to magnitudes,determine whether the systematicsources are likely to be associated with the LAEs, we calculated the probability that the NICMOS detections in the 15H field are from foreground galaxies."," In order to determine whether the sources are likely to be associated with the LAEs, we calculated the probability that the NICMOS detections in the 15H field are from foreground galaxies."
+ To do this we obtained the publicly available catalogs of the NICMOS Ultra-Deep]http://archive., To do this we obtained the publicly available catalogs of the NICMOS Ultra-Deep.
+"stsci.edu/prepds/udf/udfhlsp.html.. For footnotetext[5consistency with our aperture photometry of the LAEs, we used 0/66 diameter apertures."," For consistency with our aperture photometry of the LAEs, we used 6 diameter apertures."
+" To == 26.1, which is the aperture magnitude of the fainter of the two detected sources, we found cumulative number counts of ~60 arcmin™."," To = 26.1, which is the aperture magnitude of the fainter of the two detected sources, we found cumulative number counts of $\sim$ 60 $^{-2}$."
+" Therefore, there is less than a one per-cent chance of a foreground interloper falling within our positional error ellipse."," Therefore, there is less than a one per-cent chance of a foreground interloper falling within our positional error ellipse."
+" For the brighter source, this probability is about lower."," For the brighter source, this probability is about lower."
+ We conclude with than 99% confidence that each of the NICMOS detected greatersources is indeed an LAE discovered by Martinetal.(2008)., We conclude with greater than $99$ confidence that each of the NICMOS detected sources is indeed an LAE discovered by \cite{Martin08}.
+". Strong eemission can in some cases contribute appreciably to broadband flux density (Schaerer&deBarros2010),, and for our objects this could happen."," Strong emission can in some cases contribute appreciably to broadband flux density \citep{SdB10}, and for our objects this could happen."
+" At z25.7, iis included in the wide bbandpass, which covers approximately 0.8 to 1.4µπι."," At $z=5.7$, is included in the wide bandpass, which covers approximately 0.8 to 1.4."
+". However, the contribution is negligible for our objects- our measured ffluxes account for only of the fflux densities for the two detected sources."," However, the contribution is negligible for our objects– our measured fluxes account for only of the flux densities for the two detected sources."
+" This enhancement lies within our uncertainties, so no corrections are made to our photometry and equivalent width measurements."," This enhancement lies within our uncertainties, so no corrections are made to our photometry and equivalent width measurements."
+ These NICMOS observations provide the first direct images of two sources (MSDM 29.545 and MSDM 71-5) which had previously only been detected in our spectroscopic data., These NICMOS observations provide the first direct images of two sources (MSDM 29.5+5 and MSDM 71-5) which had previously only been detected in our spectroscopic data.
+" On the other hand, the source which is undetected inJ110., MSDM 80-23, shows a strong detection in Subaru narrowband imaging with the NB816 filter (object #555 in Murayamaetal.2007))."," On the other hand, the source which is undetected in, MSDM 80+3, shows a strong detection in Subaru narrowband imaging with the NB816 filter (object 55 in \citealt{Murayama}) )."
+" In Figure 2 we show the narrowband image of MSDM 80+3, with our position error ellipse overlaid."," In Figure \ref{stamp80+3} we show the narrowband image of MSDM 80+3, with our position error ellipse overlaid."
+ Our improved coordinates for MSDM 80+3 are roughly consistent with the NB816 source in Murayama et al., Our improved coordinates for MSDM 80+3 are roughly consistent with the NB816 source in Murayama et al.
+" The newly derived position is 0:77 of the NB816 detection, while the previous position given in Martin et al."," The newly derived position is 7 of the NB816 detection, while the previous position given in Martin et al."
+ waseast., was.
+". This position offset along the slit is still large compared to other sources, including those discussed in §2.2 and for 52 lower redshift line emitters that are in common between Martin et al."," This position offset along the slit is still large compared to other sources, including those discussed in 2.2 and for 52 lower redshift line emitters that are in common between Martin et al."
+ and the COSMOS NB816 sample., and the COSMOS NB816 sample.
+" Given the slight offset perpendicular to the slit (0744) shown in Figure 2,, one possibility is that the eemission from this object is extended, and our slit subtended an outlying region of the galaxy."," Given the slight offset perpendicular to the slit 4) shown in Figure \ref{stamp80+3}, one possibility is that the emission from this object is extended, and our slit subtended an outlying region of the galaxy."
+" Some evidence for this hypothesis is present in the Subaru narrowband image (Figure 2)), where MSDM 80-43 is marginally resolved."," Some evidence for this hypothesis is present in the Subaru narrowband image (Figure \ref{stamp80+3}) ), where MSDM 80+3 is marginally resolved."
+" The absence of continuum emission could indicate that the source is extended in the continuum, although extended eemission need not imply a similarly extended stellar continuum."," The absence of continuum emission could indicate that the source is extended in the continuum, although extended emission need not imply a similarly extended stellar continuum."
+" The upper bound of i- dropout LBG sizes is ry;« 0722 (in the rest-frame UV; Bouwensetal. 2004)), which corresponds to 5.3 pixels in our drizzled NIC2 images."," The upper bound of $i-$ dropout LBG sizes is $r_{hl} < $ 2 (in the rest-frame UV; \citealt{Bouwens04}) ), which corresponds to 5.3 pixels in our drizzled NIC2 images."
+" Using the IRAF artdata package, we verified that a galaxy of this size would be undetected in our NICMOS data at >26.5."," Using the IRAF artdata package, we verified that a galaxy of this size would be undetected in our NICMOS data at $> 26.5$."
+ It is tempting to conclude from the low-surface brightness signal at the position of MSDM 80-43 in our NICMOS image that it may be more extended in the frame UV., It is tempting to conclude from the low-surface brightness signal at the position of MSDM 80+3 in our NICMOS image that it may be more extended in the rest-frame UV.
+" However, this LAE fell on the bad central in these observations, and similar artifacts appear at other locations along the column."," However, this LAE fell on the bad central in these observations, and similar artifacts appear at other locations along the column."
+" The weight maps produced by drizzle indicate that noise is ~30% higher at these locations relative to the rest of the image, so the extended feature is not significant."," The weight maps produced by drizzle indicate that noise is $\sim 30$ higher at these locations relative to the rest of the image, so the extended feature is not significant."
+" Nevertheless, the nnon-detection could imply that MSDM 80-43 is more extended in the continuum than the other two LAEs (MSDM 29.5--5 and 71-5), which have half-light-radii (measured with SExtractor) of 07115 and 0""114."," Nevertheless, the non-detection could imply that MSDM 80+3 is more extended in the continuum than the other two LAEs (MSDM 29.5+5 and 71-5), which have half-light-radii (measured with SExtractor) of 15 and 14."
+" The hypothesis that the eemission be extended is also supported by our line flux measurement, maywhich is a factor of two lower than the Subaru"," The hypothesis that the emission may be extended is also supported by our line flux measurement, which is a factor of two lower than the Subaru"
+"operation of a turbo-molecur vacuum pump, connected to the detector cryostat through two metres of flexible vacuum pipe, did not cause any detectable vibrations.","operation of a turbo-molecur vacuum pump, connected to the detector cryostat through two metres of flexible vacuum pipe, did not cause any detectable vibrations."
+ The long-term stability during the first year of operation is compromised by the continuous adjustment and optimisation of the acclimatisation control system., The long-term stability during the first year of operation is compromised by the continuous adjustment and optimisation of the acclimatisation control system.
+" In our normal mode of operation, we use only wavelength calibration frames obtained in the evening and morning of an observation night to calibrate a given science frame in wavelength."," In our normal mode of operation, we use only wavelength calibration frames obtained in the evening and morning of an observation night to calibrate a given science frame in wavelength."
+ Nightly we obtain spectra of IAU velocity standards., Nightly we obtain spectra of IAU velocity standards.
+" Despite the variable environmental conditions, but using the pressure calibration of Sect. ??,,"," Despite the variable environmental conditions, but using the pressure calibration of Sect. \ref{sect:env_control},"
+ a standard deviation of mm/s is found using all 50 exposures of the IAU standard 1164922 in between 2009-07-09 and 2010-06-28., a standard deviation of m/s is found using all 50 exposures of the IAU standard 164922 in between 2009-07-09 and 2010-06-28.
+ Scattered light is usually an important issue in an echelle spectrograph., Scattered light is usually an important issue in an echelle spectrograph.
+" Grating efficiency is typically60-70%,, and a large part of the lost flux ends up scattered throughout the instrument."," Grating efficiency is typically, and a large part of the lost flux ends up scattered throughout the instrument."
+" Nevertheless, careful design that includes efficient baffling and use of high-quality optical components, allow HERMES to deliver very clean spectra, comparing favourably with many other echelle spectrographs."," Nevertheless, careful design that includes efficient baffling and use of high-quality optical components, allow HERMES to deliver very clean spectra, comparing favourably with many other echelle spectrographs."
+" As can be appreciated from Fig. 22,,"," As can be appreciated from Fig. \ref{fig:background},"
+ the distribution of scattered light is very local., the distribution of scattered light is very local.
+ The inter-order background signal fluctuates around of the flux in the adjacent orders., The inter-order background signal fluctuates around of the flux in the adjacent orders.
+ Only when the signal level is very low (particularly at short wavelengths) does a global straylight component become relevant., Only when the signal level is very low (particularly at short wavelengths) does a global straylight component become relevant.
+ HERMES spectra are virtually free of ghost images., HERMES spectra are virtually free of ghost images.
+" In spectra of the emission line star P Cygni with, amongst others, a strongly saturated Ha line, it turned out to be impossible to detect the presence of ghost images of the emission lines."," In spectra of the emission line star P Cygni with, amongst others, a strongly saturated $\alpha$ line, it turned out to be impossible to detect the presence of ghost images of the emission lines."
+" We presented in this contribution the design, manufacturing, and performance analyses of HERMES, a high-resolution fibre-fed spectrograph project."," We presented in this contribution the design, manufacturing, and performance analyses of HERMES, a high-resolution fibre-fed spectrograph project."
+" We showed that the instrument was built according to the demanding requirements, which resulted in an instrument of excellent stability and unsurpassed efficiency."," We showed that the instrument was built according to the demanding requirements, which resulted in an instrument of excellent stability and unsurpassed efficiency."
+ The main characteristics of our instrument are given in Table 2.., The main characteristics of our instrument are given in Table \ref{tab:summary}.
+ The ZEMAX description of the optimised optical system canbe found in Appendix ??.., The ZEMAX description of the optimised optical system canbe found in Appendix \ref{sect:zemax}. .
+being NGC 4151 (e.g. Schurch Warwick 2002. de Rosa et al.,"being NGC 4151 (e.g. Schurch Warwick 2002, de Rosa et al."
+ 2007). NGC 7582 (Bianchi et al.," 2007), NGC 7582 (Bianchi et al."
+ 2009) and NGC 1365 (Risaliti et al., 2009) and NGC 1365 (Risaliti et al.
+ 2005. 2009).," 2005, 2009)."
+ In the last two sources. in particular. large variations of the column density of the (cold) absorber occur on time scales as short as less than a day.," In the last two sources, in particular, large variations of the column density of the (cold) absorber occur on time scales as short as less than a day."
+ Here we observe also a strong variation of the ionization state of the absorbers., Here we observe also a strong variation of the ionization state of the absorbers.
+ As the primary continuum (at least the hard one) stayed almost constant. this indicates a variation of the location and/or of the density of the clouds.," As the primary continuum (at least the hard one) stayed almost constant, this indicates a variation of the location and/or of the density of the clouds."
+ It is interesting to note that Mrk 704 1s a “polar-seattering” Seyfert 1. Le. a source with optical polarization aligned perpendicularly to the radio source axis. as usually found in Seyfert 2s (Smith et al..," It is interesting to note that Mrk 704 is a “polar-scattering” Seyfert 1, i.e. a source with optical polarization aligned perpendicularly to the radio source axis, as usually found in Seyfert 2s (Smith et al.,"
+ 2004)., 2004).
+ Smith et al., Smith et al.
+ suggest that in these sources the nucleus ts seen through the edge of the torus., suggest that in these sources the nucleus is seen through the edge of the torus.
+ We fitted the spectra with two absorbing regions. either fully or partially covering the primary emission.," We fitted the spectra with two absorbing regions, either fully or partially covering the primary emission."
+ The improvement in the quality of the fit when letting the aborbers be partial is only moderate for the first observation. but more significant in the second one.," The improvement in the quality of the fit when letting the aborbers be partial is only moderate for the first observation, but more significant in the second one."
+ As the results are somewhat different in the two cases. let us discuss them separately.," As the results are somewhat different in the two cases, let us discuss them separately."
+ With the full covering absorbers. both absorbing zones are found more ionized and with a lower column density in the second observation (even if the difference in the column density for the low ionization zone is only marginal).," With the full covering absorbers, both absorbing zones are found more ionized and with a lower column density in the second observation (even if the difference in the column density for the low ionization zone is only marginal)."
+ The RGS analysis of the first observation suggests a possible (but, The RGS analysis of the first observation suggests a possible (but
+than in 1996.,than in 1996.
+ The source traces a clear pattern in the HID., The source traces a clear pattern in the HID.
+ Radio observations with the VLA were reported in Rupen. Dhawan Mioduszewski (C20052.b.eaLe.f.g.h).," Radio observations with the VLA were reported in Rupen, Dhawan Mioduszewski (2005a,b,c,d,e,f,g,h)."
+" The relatively weak radio peak occurred during a phase of X-ray flaring during the hard soft transition,", The relatively weak radio peak occurred during a phase of X-ray flaring during the hard $\rightarrow$ soft transition.
+ The recurrent black hole transient 4U 1630-47 has undergone tive outbursts between 1996 and 2005., The recurrent black hole transient 4U 1630-47 has undergone five outbursts between 1996 and 2005.
+ The source has been identified with a weak radio counterpart (Hjellming et al., The source has been identified with a weak radio counterpart (Hjellming et al.
+ 1999) but is not always detected (Gallo et al., 1999) but is not always detected (Gallo et al.
+ 2006)., 2006).
+ In Fig | we present a HID for the 1998 outburst of the source. which has the best radio coverage (Hjellming et al.," In Fig 1 we present a HID for the 1998 outburst of the source, which has the best radio coverage (Hjellming et al."
+ 1999)., 1999).
+ As with most sources. a jet has not been directly resolved in this source. so we can only rely upon the peak flux as an approximate indicator an ejection event.," As with most sources, a jet has not been directly resolved in this source, so we can only rely upon the peak flux as an approximate indicator of an ejection event."
+ This peak flux oceurs about half way. in terms of ofrelative X-ray spectral hardness. between the canonical hard and soft tor thermal dominant) X-ray states. around MJD 50861.," This peak flux occurs about half way, in terms of relative X-ray spectral hardness, between the canonical hard and soft (or thermal dominant) X-ray states, around MJD 50861."
+ By back-extrapolating the radio light curve. Hjellming et al. (," By back-extrapolating the radio light curve, Hjellming et al. ("
+1999) estimated a probably ejection date of .MID 530851. ten days earlier than this flare.,"1999) estimated a probably ejection date of $\sim$ MJD 50851, ten days earlier than this flare."
+ This source underwent a bright and prolonged X-ray outburst with complex radio behaviour in 1998., This source underwent a bright and prolonged X-ray outburst with complex radio behaviour in 1998.
+ Analysis of the radio flux anc images has recently been undertaken in Brocksopp et al. , Analysis of the radio flux and images has recently been undertaken in Brocksopp et al. (
+2007) anc Miller-Jones et al. (,2007) and Miller-Jones et al. (
+prep).,prep).
+ The peak radio flux is clearly during the hard = soft transition. but radio emission was essentially detected every time the source was observed throughout the entire period of the bright X-ray outburst.," The peak radio flux is clearly during the hard $\rightarrow$ soft transition, but radio emission was essentially detected every time the source was observed throughout the entire period of the bright X-ray outburst."
+ Observations with the VLA (Hjellming. Rupen Mioduszewski 1998: Hjellming et al.," Observations with the VLA (Hjellming, Rupen Mioduszewski 1998; Hjellming et al."
+ 1998b: Miller-Jones et al., 1998b; Miller-Jones et al.
+ in prep) clearly show that a lot of this is due to resolved component which appear to interact with the environment and/or each other., in prep) clearly show that a lot of this is due to resolved component which appear to interact with the environment and/or each other.
+ Note that extrapolating the early proper motions of the resolved maps in Miller-Jones et al. G, Note that extrapolating the early proper motions of the resolved maps in Miller-Jones et al. (
+in prep) indicates that the firs ejection every took place before the first X-ray observations. which may in principle have still been in the canonical low/hard state.,"in prep) indicates that the first ejection every took place before the first X-ray observations, which may in principle have still been in the canonical low/hard state."
+ The equally plausible alternative is that the initial proper motions were much higher. and the ejecta decelerated as they interacted with the ISM tas in e.g. XTE JISS0-S64 — Corbel et al.," The equally plausible alternative is that the initial proper motions were much higher, and the ejecta decelerated as they interacted with the ISM (as in e.g. XTE J1550-564 – Corbel et al."
+ 2001)., 2001).
+ XTE J20124381 is a black hole candidate X-ray transient that went into outburst in 1998 (e.g. Campana et al., XTE J2012+381 is a black hole candidate X-ray transient that went into outburst in 1998 (e.g. Campana et al.
+ 2002)., 2002).
+ The RXTE coverage is sparse., The RXTE coverage is sparse.
+ Hjellming. Rupen Mioduszewski (19986) and Pooley (1998) report on a weak but variable radio counterpart.," Hjellming, Rupen Mioduszewski (1998c) and Pooley (1998) report on a weak but variable radio counterpart."
+ The variability reported by Pooley (1998) may indicate an ejection around May 30 (MJD 50963) XTE 11550-3564 is a fascinating source with a wealth of interesting properties in terms of its disc-jet coupling. including powerful. moving. X-ray jets (Corbel et al.," The variability reported by Pooley (1998) may indicate an ejection around May 30 (MJD 50963) XTE J1550-564 is a fascinating source with a wealth of interesting properties in terms of its disc–jet coupling, including powerful, moving, X-ray jets (Corbel et al."
+ 2002) and large loops in the HID (Fig |»., 2002) and large loops in the HID (Fig 1).
+ Hannikainen et al. (, Hannikainen et al. (
+2001) report VLBI images of relativistic jets following a major flaring event around MID 51078. with strong evidence for a secondary core ejection event around MID 51080: Homan et al. (,"2001) report VLBI images of relativistic jets following a major flaring event around MJD 51078, with strong evidence for a secondary core ejection event around MJD 51080; Homan et al. ("
+2001) report a bright. optically thin radio source around MID 51248.,"2001) report a bright, optically thin radio source around MJD 51248."
+ These dates are indicated in Fig |., These dates are indicated in Fig 1.
+ Note that there several unpublished radio observations of this outburst which may in the future shed some further light on the evolution of this outburst., Note that there several unpublished radio observations of this outburst which may in the future shed some further light on the evolution of this outburst.
+ XTE J18594+226 is a source for which there was both good X-ray and radio coverage around the brightest phases of the outburst (see e.g. Casella et al., XTE J1859+226 is a source for which there was both good X-ray and radio coverage around the brightest phases of the outburst (see e.g. Casella et al.
+ 2004 for an X-ray study of the outburst)., 2004 for an X-ray study of the outburst).
+ Radio data for this bright outburst are from Brocksopp et al. (, Radio data for this bright outburst are from Brocksopp et al. (
+2002). who noted five radio flare events in a sequence of declining strength (reminiscent of those seen in GRS 19154105 e.g. Fender et al.,"2002), who noted five radio flare events in a sequence of declining strength (reminiscent of those seen in GRS 1915+105 e.g. Fender et al."
+ 1999)., 1999).
+ All five radio flare events occur during a prolonged period of X-ray flaring around the middle of the hard to soft transition. with the strongest of these happening at the time of a local peak in the intensity.," All five radio flare events occur during a prolonged period of X-ray flaring around the middle of the hard to soft transition, with the strongest of these happening at the time of a local peak in the X-ray intensity."
+ Fig 2 presents a close-up of the HID around the time of these five flare events., Fig 2 presents a close-up of the HID around the time of these five flare events.
+ All of these events occur around the same hardness. with the greatest outlier being revealed as associated with a brief and rapid hardening of the X-ray spectrum of total duration XOLS days.," All of these events occur around the same hardness, with the greatest outlier being revealed as associated with a brief and rapid hardening of the X-ray spectrum of total duration $\leq 1.8$ days."
+ This is explored in more detail in the discussion (see Fig 2)., This is explored in more detail in the discussion (see Fig 2).
+ Corbel et al. (, Corbel et al. (
+2004) present a detailed discussion of the ATCA radio observations of this source during its 2001 outburst.,2004) present a detailed discussion of the ATCA radio observations of this source during its 2001 outburst.
+ The peak radio flux is detected close to the start of the hard = soft transition. although there is subsequently a lack of radio observations for 15 days and the true peak may have been missed.," The peak radio flux is detected close to the start of the hard $\rightarrow$ soft transition, although there is subsequently a lack of radio observations for 15 days and the true peak may have been missed."
+ Subsequently the radio emission is found to be temporarily undetectable. but then to re-emerge in the soft state.," Subsequently the radio emission is found to be temporarily undetectable, but then to re-emerge in the soft state."
+ Park et al. (, Park et al. (
+2004) and Kalemei et al. (,2004) and Kalemci et al. (
+2005). report in. detail multiwavelength observations of this recurrent. transient.,2005) report in detail multiwavelength observations of this recurrent transient.
+ In particular. Park et al. (," In particular, Park et al. ("
+2004) detected a rapidly rising radio flare on MJD 52443.,2004) detected a rapidly rising radio flare on MJD 52443.
+ During the decay of the outburst. radio emission is observed to reactivate by MJD 52487. a week after the MJD 52480 date that Kalemei et al. (," During the decay of the outburst, radio emission is observed to reactivate by MJD 52487, a week after the MJD 52480 date that Kalemci et al. ("
+2005) report for the soft to hard state transition. based on timing as well as spectral properties.,"2005) report for the soft to hard state transition, based on timing as well as spectral properties."
+ The 2003 outburst of this recurrent transient received excellent radio and X-ray coverage. as reported in McClintock et al. (," The 2003 outburst of this recurrent transient received excellent radio and X-ray coverage, as reported in McClintock et al. ("
+2006).,2006).
+ The HID reveals a peak in the radio emission near the start (i.e. hard side) of the hard to soft transition (as in XTE 118594226 this is associated with a phase of X-ray flaring). with many more detections in the intermediate and initial soft states before a large number of upper limits in the later soft state and during the soft to hard transition.," The HID reveals a peak in the radio emission near the start (i.e. hard side) of the hard to soft transition (as in XTE J1859+226 this is associated with a phase of X-ray flaring), with many more detections in the intermediate and initial soft states before a large number of upper limits in the later soft state and during the soft to hard transition."
+ Brocksopp et al. (, Brocksopp et al. (
+2005) report radio coverage of this outburst.,2005) report radio coverage of this outburst.
+ The peak observed radio emission occurs around MJD 52655. at which point the source is already in a rather soft state.," The peak observed radio emission occurs around MJD 52655, at which point the source is already in a rather soft state."
+ Subsequently there are a number of non-detections in the soft state., Subsequently there are a number of non-detections in the soft state.
+ Brocksopp et al. (, Brocksopp et al. (
+2005) also report that the radio emission in this source appears to switeh back on between MID 32715-52728 (increasing by a,2005) also report that the radio emission in this source appears to switch back on between MJD 52715–52728 (increasing by a
+in eeneral: (he principal axis of inertia ο (lhe dipole axis of the natal magnetic field) is inclined. al an angle a to the rotation axis.,in general; the principal axis of inertia ${\bf e}_3$ the dipole axis of the natal magnetic field) is inclined at an angle $\alpha$ to the rotation axis.
+ Gravitational waves are emitted by the deformed star at the spin frequency Q=//2 and its first harmonic 2€?=f (unless a= 7/2. when there is no //2 component). and the spin-down Iuminosity is proportional to (16sin?a+cos?)sin? (??)..," Gravitational waves are emitted by the deformed star at the spin frequency $\Omega=f/2$ and its first harmonic $2\Omega=f$ (unless $\alpha=\pi/2$ , when there is no $f/2$ component), and the spin-down luminosity is proportional to $(16\sin^2\alpha + \cos^2\alpha) \sin^2\alpha$ \citep{sha83,bon96}."
+ Upon averaging over a. polarization. and orientation (the position angle of the rotation axis on the skv cannot normally be measured). one can define the characteristic gravitational wave strain A.=(1282!/15)!?GL.[?e/(de?) (?).. which recuces (ο upon substituting (8)).," Upon averaging over $\alpha$, polarization, and orientation (the position angle of the rotation axis on the sky cannot normally be measured), one can define the characteristic gravitational wave strain $h_{\rm c} = (128\pi^4 /15)^{1/2} G I_{zz} f^2 \epsilon/ (dc^4)$ \citep{bra98}, which reduces to upon substituting \ref{eq:gra6d}) )."
+" Polar magnetic burialtherefore generates eravilational radiation whose amplitude fh.zz6x10.79 (for tvpical parameters M,ZM.~10OAL, and b=30) is c10? limes greater (han that produced by the natal. undistorted magnetic dipole. 5,zz (022).. due to the enhanced. Maxwell stress. [rom the compressed equatorial magnetic field."," Polar magnetic burialtherefore generates gravitational radiation whose amplitude $h_{\rm c} \approx 6\times 10^{-26}$ (for typical parameters $M_{\rm a}\gtrsim M_{\rm c} \sim 10^{-5} M_\ast$ and $b=30$ ) is $\sim 10^5$ times greater than that produced by the natal, undistorted magnetic dipole, $h_{\rm c} \approx 10^{-31} 
+ (B/10^{12}\,{\rm G})^2 (f/0.6\,{\rm kHz})^2
+ (d/1\,{\rm kpc})^{-1}$ \citep{kat89,bon96}, , due to the enhanced Maxwell stress from the compressed equatorial magnetic field."
+ The self-consistentform of Fe) defined by (3)) is needed to caleulate this stress properly: citet melü1.., The self-consistentform of $F(\psi)$ defined by \ref{eq:gra4}) ) is needed to calculate this stress properly; \\citet{mel01}.
+ Figure Saisaplolofh. versus f for 10.<MM.10? and 10€b<10°., Figure \ref{fig:gra3}$ $a$ is a plot of $h_{\rm c}$ versus $f$ for $10^{-8} \leq M_{\rm a}/M_{\sun} \leq 10^{-2}$ and $10\leq b\leq 10^2$.
+ The sensitivity curves for initial LIGO (L) and advanced LIGO (AL) are superimposed. corresponding io the weakest source detectable with 99 per cent confidence in 7=10*s of integration time. if the frequency. anc pliase of the signal at the detector are known in advance (??)..," The sensitivity curves for initial LIGO $L$ ) and advanced LIGO $AL$ ) are superimposed, corresponding to the weakest source detectable with 99 per cent confidence in $\tau=10^7\,{\rm s}$ of integration time, if the frequency and phase of the signal at the detector are known in advance \citep{bra98,sch99}."
+" The plot suggests that the prospects of detecting objects with M,=10M. are encouraging. a point first made in the context. of magnetic burial by 2.."," The plot suggests that the prospects of detecting objects with $M_{\rm a} \gtrsim 10^{-5} M_{\sun}$ are encouraging, a point first made in the context of magnetic burial by \citet{mel00}."
+ Such objects include the five accreting. X-ray. millisecond. pulsars discovered. at the time of writing. SAX JI1308.4—3653.," Such objects include the five accreting X-ray millisecond pulsars discovered at the time of writing, SAX $-$ 3658, XTE $-$ 305, XTE $-$ 314, XTE $-$ 294, and XTE $-$ 338 \citep{wij98,cha98,gal02,mar02,str03,cam03}."
+ is facilitated if the time-dependent Doppler shift from (the binary orbit is known well enough to be subtracted. but this is only practical for some objects of known /. not for an all-sky search (?)..," Detectability is facilitated if the time-dependent Doppler shift from the binary orbit is known well enough to be subtracted, but this is only practical for some objects of known $f$, not for an all-sky search \citep{bra98}."
+ The characteristic gravitational wave strain ὃν is a lower limit in an important sense: it contains the assumption (hat. after averaging over polarization and orientation. (he bulk of the gravitational wave signal is emitted al f=20.," The characteristic gravitational wave strain $h_{\rm c}$ is a lower limit in an important sense: it contains the assumption that, after averaging over polarization and orientation, the bulk of the gravitational wave signal is emitted at $f=2\Omega$."
+" This is true when the observer views the system along its rotation axis and measures oxAhxe?""sin?a in the two polarizations.", This is true when the observer views the system along its rotation axis and measures $h^+ \propto h^\times \propto e^{2i\Omega t} \sin^2\alpha$ in the two polarizations.
+" It is nol true when the observer views (he svstem perpendicular (to ils rotationaxis ancl measures frxe""YO,si«2a and h*xe?""hEsina cosa. lor; example 22 (2?).."," It is true when the observer views the system perpendicular to its rotationaxis and measures $h^+ \propto e^{2i\Omega t} \sin^2\alpha$ and $h^\times \propto e^{i\Omega t} \sin\alpha\cos\alpha$ , for example \citep{sha83,bon96}. ."
+ One certainly1 expects. as a matter of chance. to observe some svstems in the latter orientation. or eloseto it.," One certainly expects, as a matter of chance, to observe some systems in the latter orientation, or closeto it."
+For these sources. (he average quantitv ὃς substantially underestimates (he (rue wave strain,"For these sources, the average quantity $h_{\rm c}$ substantially underestimates the true wave strain"
+Iu the standard flare scenario. electrons that were accelerated in the corona stream dowmwards toward the deuse lavers of the solar atmosphere. where they are stopped via collisions producing intense hard X-rav (IIXBR) cimission dn the chromosphere.,"In the standard flare scenario, electrons that were accelerated in the corona stream downwards toward the dense layers of the solar atmosphere, where they are stopped via collisions producing intense hard X-ray (HXR) emission in the chromosphere."
+ Hüeher energv electrous penetrate deepero iuto the chromosphere., Higher energy electrons penetrate deeper into the chromosphere.
+ Therefore. imieasurenieuts of the spatial structure of the TINR cunission as a function of energv provide information about the chromospheric density.," Therefore, measurements of the spatial structure of the HXR emission as a function of energy provide information about the chromospheric density."
+ Όσιο optically thin. N-ravs give the most direct iuformation about the spatial aud. cucrev distribution of energetic electrous in the solar atmosphere.," Being optically thin, X-rays give the most direct information about the spatial and energy distribution of energetic electrons in the solar atmosphere."
+ Prior to the launch of the Ramaty ]Hiüeh Encrey Solar Spectroscopic huager (RITESSI) (2).. INR instrmmcuts typically had Iuited dmagine-spectroscopy capabilities (?)..," Prior to the launch of the Ramaty High Energy Solar Spectroscopic Imager (RHESSI) \citep{lin2002}, HXR instruments typically had limited imaging-spectroscopy capabilities \citep{Kosugi1992yohkoh}."
+ RIIESSIES aptitude to image m various energy ranges opens new horizous for studviusg the detailed structure of IINR cluitting sources iu the chromosphere., RHESSI's aptitude to image in various energy ranges opens new horizons for studying the detailed structure of HXR emitting sources in the chromosphere.
+ RIIESSI does uot directly image the Sun but uses 9 pairs of Rotating Modulation Collimators (RAICS) to time-modulate spatial information in the signal obtained in ifs eecrmnaniun detectors (?7)., RHESSI does not directly image the Sun but uses 9 pairs of Rotating Modulation Collimators (RMCs) to time-modulate spatial information in the signal obtained in its germanium detectors \citep{hurford2002}.
+ Each RAIC has a different thickuess of its slits aud slats makine it sensitive to differeut spatial scales. providing modulation at nine spatial frequencies.," Each RMC has a different thickness of its slits and slats making it sensitive to different spatial scales, providing modulation at nine spatial frequencies."
+ The reconstruction of au nuage frou these time modulated lightcurves. can be accomplished by various nuage algorithius (??7777)..," The reconstruction of an image from these time modulated lightcurves, can be accomplished by various imaging algorithms \cite[][]{emslie2003,BattagliaBenz2007,KruckerLin2008,Saint-Hilaire_etal2008,dennis2009}."
+ The now visibility based approach to RITESSI imagine starts by πιο (stacking) the lehtcurves per roll bins over a few spin periods of the spacecraft (?).., The new visibility based approach to RHESSI imaging starts by summing (stacking) the lightcurves per roll bins over a few spin periods of the spacecraft \citep{schmahl2007}.
+ The fitted amplitudes aud the plases in the individual roll bius are X-ray visibilities., The fitted amplitudes and the phases in the individual roll bins are X-ray visibilities.
+ This effectively provides two dinuensional spatial Fourier componcuts (X-ray visibilitices} over a wide range of energies (77).," This effectively provides two dimensional spatial Fourier components (X-ray visibilities) over a wide range of energies \citep{hurford2002,schmahl2007}."
+. To convert the time-moculated signal or X-ray visibilitics to an image is an inverse problem (e.g.??)..," To convert the time-modulated signal or X-ray visibilities to an image is an inverse problem \citep[e.g.][]{Piana_etal2007,Prato_etal2009}."
+" The reconstructed nuages face unavoidable difficultics due to nieasuremient errors, finite coverage in Fourier space. and ill-poseduess of the reconstruction problem."," The reconstructed images face unavoidable difficulties due to measurement errors, finite coverage in Fourier space, and ill-posedness of the reconstruction problem."
+ The resulting reconstruction errors and snall dynamic rauge lakes it difficult to accurately measure source sizes fron reconstructed images., The resulting reconstruction errors and small dynamic range makes it difficult to accurately measure source sizes from reconstructed images.
+ At best the imaging resolution is down to 2 arcseconds but im practice is around 7 arcseconds for the typical flare nouthermal energy range., At best the imaging resolution is down to 2 arcseconds but in practice is around 7 arcseconds for the typical flare nonthermal energy range.
+ However. the moments of N-rav source distribution (source position. source size. etc} can be inferred with higher precision either from the time-modulated signal (7777) or visibilitics (2777)..," However, the moments of X-ray source distribution (source position, source size, etc) can be inferred with higher precision either from the time-modulated signal \cite[][]{aschwanden2002,KruckerLin2008,Saint-Hilaire_etal2008,Fivian_etal2009}
+ or visibilities \citep{Xu_etal2008,Kontar_etal08,dennis2009,Prato_etal2009}."
+ Thus. RITESSI nieasurenients of N-ray source positions can recover sub-arcsecond information using RIIESSI modulated liehteurves (227). or visibilities (227)...," Thus, RHESSI measurements of X-ray source positions can recover sub-arcsecond information using RHESSI modulated lightcurves \citep{aschwanden2002,Liu_etal2006,Mrozek06}
+ or visibilities \citep{Kontar_etal08,dennis2009,Prato_etal2009}."
+ This has allowed clear demonstration of the heielt-cncrey dependence of IIXR sources: higher euergv sources originate at lower heights (???)..," This has allowed clear demonstration of the height-energy dependence of HXR sources: higher energy sources originate at lower heights \citep{aschwanden2002,brown2002,Liu_etal2006}."
+ This has substautially improved upon previous results with Yohkoh/IENT (?).., This has substantially improved upon previous results with Yohkoh/HXT \citep{1992PASJ...44L..89M}.
+ The recently developed visibility-based technique (2).. allowed ? to iuprove previous nieasureiments and infer characteristic sizes (FWIIM) of INR footpoiuts and not just the ceutroid height.," The recently developed visibility-based technique \citep{schmahl2007}, allowed \citet{Kontar_etal08} to improve previous measurements and infer characteristic sizes (FWHM) of HXR footpoints and not just the centroid height."
+ From this the convergence of the magnetic flux aud neutral hydrogen deusitv distribution in the chromosphere was inferred., From this the convergence of the magnetic flux and neutral hydrogen density distribution in the chromosphere was inferred.
+ Tn this paper we study the structure of solar flare hard X-ray sources using various dmaecine aleoritlius: PINON. MEN. CLEAN aud visibility forward fit.," In this paper we study the structure of solar flare hard X-ray sources using various imaging algorithms: PIXON, MEM, CLEAN and visibility forward fit."
+ Using RIIESSI X-rav visibilities we fud the characteristic shapes aud positions for differcut energy rauges., Using RHESSI X-ray visibilities we find the characteristic shapes and positions for different energy ranges.
+ We show that the technique of forward fitting X-ray visibilitics allows us to determine not oulv the FEWIIMD of the, We show that the technique of forward fitting X-ray visibilities allows us to determine not only the FWHM of the
+The shocked wind aud shocked ejecta are iu pressure baluice aud separated by the contact discontimuitv.,The shocked wind and shocked ejecta are in pressure balance and separated by the contact discontinuity.
+" The ejecta ect decelerated at Ris2:5«1019E,sy(Ty/2)oy,2M.5lege. where Ej aud Ey are. respectively.. he total kinetic cnerey and the initial Lorentz factor of the semiaclativistic ejecta."," The semi-relativistic ejecta get decelerated at $R_{\rm HN}\simeq 5\times10^{16} E_{k,51}
+(\Gamma_0/2)^{-2}\dot{M}_{-5}^{-1} v_{w,3} {\rm cm}$ , where $E_{k}$ and $\Gamma_0$ are, respectively, the total kinetic energy and the initial Lorentz factor of the semi-relativistic ejecta."
+" Since the shock expausion is expecος to be adiabatic. we have where [μιςcL/SPT?Pp,(Rye? (Blandford&Melxoe1976) is the total cuereywave. LagmD|peByzμια IS the total euergy iu the shocked ejecta. Τ is the Lore""tz factor of the shocked material aud An is the wind D at radius R."," Since the shock expansion is expected to be adiabatic, we have where $E_{\rm sh,w}\simeq
+4/3\pi\sigma\beta^2\Gamma^2R^3\rho_w(R)c^2$ \citep{BM76} is the total energy, $E_{\rm sh,ej}\simeq \Gamma^2(e'+\rho'c^2+\beta^2p')V_{\rm sh,ej}$ is the total energy in the shocked ejecta, $\Gamma$ is the Lorentz factor of the shocked material and $\rho_w(R)$ is the wind density at radius $R$."
+" Hore 0 is a cocfBicieut with o»0,35 wwjon >j banda>O73 when >O.", Here $\sigma$ is a coefficient with $\sigma\rightarrow 0.35$ when $\beta\rightarrow 1$ and $\sigma\rightarrow 0.73$ when $\beta\rightarrow 0$.
+ We use c=mou0.73 to bridge the snp between the non-relativistic Case al ultra-relativistic case. as used in ," We use $\sigma=0.73-0.38\beta$ to bridge the gap between the non-relativistic case and ultra-relativistic case, as used in \citet{HDL98}."
+"IIuaugetal. 1l)pae. p=e),1 and p!=(35Le are respectively. the internal eueres> density. rest mass densitv aud pressure in the comoving frame of shocked. material. aud 5 is the adiabatic iudex."," $e'=\frac{\hat{\gamma}\gamma_{\rm
+rs}+1}{\hat{\gamma}-1}(\gamma_{\rm rs}-1)\rho_{\rm ej}c^2$ , $\rho'=\frac{\hat{\gamma}\gamma_{\rm rs}+1}{\hat{\gamma}-1}\rho_{\rm
+ej}$ and $p'=(\hat{\gamma}-1)e'$ are respectively, the internal energy density, rest mass density and pressure in the comoving frame of shocked material, and $\hat{\gamma}$ is the adiabatic index."
+ Wy.—1)/(424)1 is the volume of shocked ejecta. where Vy is the volume of the ejecta before beiug shocked.," $V_{\rm sh,ej}=V_0/(\frac{\hat{\gamma}\gamma_{\rm
+rs}+1}{\hat{\gamma}-1})$ is the volume of shocked ejecta, where $V_0$ is the volume of the ejecta before being shocked."
+ Tere 54 is the Loreutz factor of the reverse shock. which relates with the Loreutz factor D of the shocked material by 244;2PoP1ye}. Where oy and ος are the velocities of the uushockedejecta aud shiocked ejecta respectively.," Here $\gamma_{\rm rs}$ is the Lorentz factor of the reverse shock, which relates with the Lorentz factor $\Gamma$ of the shocked material by $\gamma_{\rm rs}\simeq \Gamma_0\Gamma(1-\beta_0\beta_{\rm se})$, where $\beta_0$ and $\beta_{\rm se}$ are the velocities of the unshockedejecta and shocked ejecta respectively."
+ Solving Eq. CAT)), Solving Eq. \ref{EC}) )
+ with typical values for the parameters. we finally eet the Lorentz factor and speed ofthe reverse shock.," with typical values for the parameters, we finally get the Lorentz factor and speed of the reverse shock,"
+Over the past few. vears. searches for Lvman-break galaxies (LBGs) and Lyman-alpha emitters (LAL) have revealed new populations of voung. star-forming objects at redshillt bevond z=6 (eg.Starketal.2OOTa:Luchardct2008:Bouwensetal.2008.2009) that shed light on the star formation history at high redshifts and the sources responsible for cosmic reionization.,"Over the past few years, searches for Lyman-break galaxies (LBGs) and Lyman-alpha emitters (LAE) have revealed new populations of young, star-forming objects at redshift beyond $z=6$ \citep[e.g.][]{Stark07,
+Richard08, Bouwens08a, Bouwens08b} that shed light on the star formation history at high redshifts and the sources responsible for cosmic reionization."
+ In addition. evidence as been found for a population of massive. evolved systems around z=5 (Alobasheretal.2005:Wiklind2008) hat hints at an earlier period of higher star formation than las vet to be seen in surveys of bright LDBCGs and LALs.," In addition, evidence has been found for a population of massive, evolved systems around $z=5$ \citep{Mobasher05,
+Wiklind08} that hints at an earlier period of higher star formation than has yet to be seen in surveys of bright LBGs and LAEs."
+ Theoretically. it is expected. that the bulk of. the star formation during reionization had taken place in less uminous galaxies than previously obscrved(c.g.Barkana&Loeb2001:WyitheLoch 2006).," Theoretically, it is expected that the bulk of the star formation during reionization had taken place in less luminous galaxies than previously \citep[e.g.][]{BL01, WL06}."
+. Dwarf galaxies are supposed to supply the bulk of the radiation that reionizes he Universe as well as provide the building blocks of the Milky Way andthe possible birth place of globular clusters (c.g.Ixravtsov.&Cinedin2005:WKravtsoy2006:Mooreetal.2006:Macau2008:Munoz 2009).," Dwarf galaxies are supposed to supply the bulk of the radiation that reionizes the Universe as well as provide the building blocks of the Milky Way andthe possible birth place of globular clusters \citep[e.g.][]{KG05, GK06, Moore06, Madau08b,
+Munoz09}."
+. Attempts to probe the faint end of the luminosity function aave often sacrificed [eld-of-view for the sake of higher [ux sensitivity., Attempts to probe the faint end of the luminosity function have often sacrificed field-of-view for the sake of higher flux sensitivity.
+ Objects in the Li2aremin? Hubble Ultra Deep Field (IUDE). for example. can be seen in the σκσυAp passbancl down to an apparent magnitude of about 29.," Objects in the $11.2\, \rm{arcmin}^2$ Hubble Ultra Deep Field (HUDF), for example, can be seen in the $z_{850,{\rm AB}}$ passband down to an apparent magnitude of about $29$."
+ In an even more extreme case. a single long-slit: spectroscopic survey for eravitationally-lensecl LAs has a field-of-view of only about 32aresec? (which is even smaller in the unlensed source plane) but can achieve large boosts in sensitivity if positioned on the eravitational lensing critical line of a foreground galaxy cluster.," In an even more extreme case, a single long-slit spectroscopic survey for gravitationally-lensed LAEs has a field-of-view of only about $32\, \rm{arcsec}^2$ (which is even smaller in the unlensed source plane) but can achieve large boosts in sensitivity if positioned on the gravitational lensing critical line of a foreground galaxy cluster."
+" Phe James Webb Space Telescope planned. for launch in 2013. will have a better sensitivity. enabling observers to probe ""pencil beams” out to even higher redshifts."," The James Webb Space Telescope ), planned for launch in 2013, will have a better sensitivity, enabling observers to probe “pencil beams” out to even higher redshifts."
+ Recent data analysis from. deep. anc narrow galaxy surveys has begun to be incorporated. into theoretical models of structure formation and. reionization., Recent data analysis from deep and narrow galaxy surveys has begun to be incorporated into theoretical models of structure formation and reionization.
+ Several studies have attempted: to relate the observed. UV. and Lyman-alpha bIuminosities from these sources to the masses of the dark matter halos that. host. them by. fitting results from. simulations and semi-analvtic prescriptions to the observations and so understand. the clustering ancl abundance of such galaxies (e.g.LeDelliouetal.2005.2006: 2009)..," Several studies have attempted to relate the observed UV and Lyman-alpha luminosities from these sources to the masses of the dark matter halos that host them by fitting results from simulations and semi-analytic prescriptions to the observations and so understand the clustering and abundance of such galaxies \citep[e.g.][]{LeDelliou05, LeDelliou06, SLE07, KTN07, Nagamine08, Orsi08, KTN09}. ."
+ Ln ασΠίο. it," In addition, it"
+"neutral medium ionizing aud compressing the medi, and leaving a lasting impression.","neutral medium, ionizing and compressing the medium and leaving a lasting impression."
+ Not only do rregious and SNRs impact the ISAL the structure of the ISAL - particularly density cuhancemeuts - affects the morphology of rregious and SNRs.," Not only do regions and SNRs impact the ISM, the structure of the ISM - particularly density enhancements - affects the morphology of regions and SNRs."
+ After the SNR or iregion has ceased to exist iu continu. the imprint may remain in the ISM in the form of au sshell.," After the SNR or region has ceased to exist in continuum, the imprint may remain in the ISM in the form of an shell."
+ Because of the Iarge range of spatial scales sampled with the combined Parkes and ATCA data. the SGPS is an ideal dataset iu which to explore these effects.," Because of the large range of spatial scales sampled with the combined Parkes and ATCA data, the SGPS is an ideal dataset in which to explore these effects."
+ Iu this paper we introduce the details of the Southern Galactic Plane Survey with attention to sspectral line and A21-cim continua data from the SGPS Test Region (32525xF<33385. 005xb 5.2).," In this paper we introduce the details of the Southern Galactic Plane Survey with attention to spectral line and $\lambda$ 21-cm continuum data from the SGPS Test Region $325\fdg5 \leq l \leq 333\fdg5$; $-0\fdg5 \leq b \leq
++3\fdg5$ )."
+ Other scieutific highlichts from the SCPS are discussed clsewhere: two large sshells discovered in the Parkes dataare presented in MeClure-Cuiffithsetal.(2000b).. preliminary images of eecnission aud absorption features are presented in MeClure-Giifithsetal.(2000a).. and the polarization properties of the Test Region are presented in Cacusleretal. (2000a).," Other scientific highlights from the SGPS are discussed elsewhere: two large shells discovered in the Parkes dataare presented in \citet{mcgriff00b}, preliminary images of emission and absorption features are presented in \citet{mcgriff00a}, and the polarization properties of the Test Region are presented in \citet{gaensler00a}."
+. Tere we explore the counectious between the aand A21-0 contimmiun inages of the Test Reeiou., Here we explore the connections between the and $\lambda$ 21-cm continuum images of the Test Region.
+ In rofsubseciobj 2. we describe the survey objectives. observing aud data analysis strategies.," In \\ref{subsec:obj} \ref{sec:obs} we describe the survey objectives, observing and data analysis strategies."
+ Iu rofsecicont we discuss the A21-c11 continuum emission., In \\ref{sec:cont} we discuss the $\lambda$ 21-cm continuum emission.
+ aabsorptiou towards σοι sources is discussed in rofseciabs.., absorption towards continuum sources is discussed in \\ref{sec:abs}.
+ We have cliosen a representative sample of rregious and supernova ποιαές (SNRs) to study the relationship between the coutiuuun enüssion from these objects aud the surrounding eenvironnents mn refsecienis.., We have chosen a representative sample of regions and supernova remnants (SNRs) to study the relationship between the continuum emission from these objects and the surrounding environments in \\ref{sec:emis}.
+ The Test Ποσο is au excelleut area to initiate such a study as it coutaius many rregious aud SNRs. as well as extended euissiou structure.," The Test Region is an excellent area to initiate such a study as it contains many regions and SNRs, as well as extended emission structure."
+ Usine aahsorption and nunorphological matches to the continuum ciission. we seek to create a three-dimensional view of the Calaxy iu this subregion.," Using absorption and morphological matches to the continuum emission, we seek to create a three-dimensional view of the Galaxy in this subregion."
+ The general goal of the Southern Galactic Plane Survey is to provide a dataset with which to study the structure and dynamics of the neutral lvdrogen 1)) in the inner Calaxy., The general goal of the Southern Galactic Plane Survey is to provide a dataset with which to study the structure and dynamics of the neutral hydrogen ) in the inner Galaxy.
+ Previous studies of the inner Galaxy have lacked the seusitivitv aud resolution necessary to study the physical processes of the interstellar iiediun (ISM) over a large range of spatial scales., Previous studies of the inner Galaxy have lacked the sensitivity and resolution necessary to study the physical processes of the interstellar medium (ISM) over a large range of spatial scales.
+ Though the specific eoals of the SGPS are umimerous. we will highlieht a few below: Observations of the SCPS Test Region were made with the Australia Telescope Compact Array (ATCA: Frater. Brooks. Whiteoalk 1992) and Parkes Radio Telescope.," Though the specific goals of the SGPS are numerous, we will highlight a few below: Observations of the SGPS Test Region were made with the Australia Telescope Compact Array (ATCA; Frater, Brooks, Whiteoak 1992) and Parkes Radio Telescope."
+ The ATCA is an cast-west svuthesis instrument near Narrabri NSW. with six 22 ni antennas on a 6 kin track.," The ATCA is an east-west synthesis instrument near Narrabri NSW, with six 22 m antennas on a 6 km track."
+ Five auteunas are movable iuto coufiguratious with baselines between 31 12 aud 6 kin., Five antennas are movable into configurations with baselines between 31 m and 6 km.
+ The ATCA data cousist of a 190 pointing mosaic covering 32575x4<33375 and Whxbox(395., The ATCA data consist of a 190 pointing mosaic covering $325\fdg5 \leq l \leq 333\fdg5$ and $-0\fdg5 \leq b \leq +3\fdg5$.
+ These data were obtaineddurius five separate observing sessions between April 1997 aud April L998., These data were obtainedduring five separate observing sessions between April 1997 and April 1998.
+ The observing dates and times are given in Table 1.., The observing dates and times are given in Table \ref{tab:obs}.
+" The observations were made with several compact array configurations - TOA. τος, 750D. and 375 - in order to obtain maxima sensitivity to large scale structures;"," The observations were made with several compact array configurations - 750A, 750C, 750D, and 375 - in order to obtain maximum sensitivity to large scale structures."
+ Each of the 190 poiutiugs was observed in forty 30 8 snapshots at a broad range of hour augles for goodwer coverage., Each of the 190 pointings was observed in forty 30 s snapshots at a broad range of hour angles for good coverage.
+ The pointines for the Test Reeiou were arranged onu a square evid with sseparation (Nyquist) as determined by the ATCA primary beam FWHAL which is aat. A-21 cm.," The pointings for the Test Region were arranged on a square grid with separation (Nyquist) as determined by the ATCA primary beam FWHM, which is at $\lambda$ -21 cm."
+ The pointing ceuters are plotted ou the 2l-cm coutinuuu ATCA image of the δις Test Reeion shown in Fig. l.., The pointing centers are plotted on the 21-cm continuum ATCA image of the SGPS Test Region shown in Fig. \ref{fig:centers}.
+ The ATCA feeds receive two orthogonal linear polarizations. XX aud Y.," The ATCA feeds receive two orthogonal linear polarizations, $X$ and $Y$ ."
+" All observations wererecorded iu a wideband continuum mode with 32 channels. each | MITz. across a 128 ΑΠ, total bandwidth with polarizationproducts VY. Y Y. XY. and YX to enable calculation of all four Stokes parameters: aud"," All observations wererecorded in a wideband continuum mode with 32 channels, each 4 MHz, across a 128 MHz total bandwidth with polarizationproducts $XX$ , $YY$ , $XY$ , and $YX$ to enable calculation of all four Stokes parameters; and"
+which are not associated wilh prominence eruptions.,which are not associated with prominence eruptions.
+ The residual acceleration for the very impulsve 2007 December 31 CME was 90ms7. while for the CMEs on 2007 November 16 and 2009 December 16. it was found to be Ίδης7 and —2ims7. respectively.," The residual acceleration for the very impulsive 2007 December 31 CME was $90\mpss$, while for the CMEs on 2007 November 16 and 2009 December 16, it was found to be $18\mpss$ and $-2\mpss$, respectively."
+ The other CMESs. which are associated with prominences do not show such an acceleration profile.," The other CMEs, which are associated with prominences do not show such an acceleration profile."
+ Chen&Ixrall(2003). have invoked the flix injection mechanism (o trigger an eruption in a magnetic flix rope. which leads to the residual acceleration phase.," \citet{Chen.Krall2003} have invoked the flux injection mechanism to trigger an eruption in a magnetic flux rope, which leads to the residual acceleration phase."
+ In the cases analvsecl here. we observe that only the flare-associated CMES undergo residual acceleration. which indicates (hat Εικ injection seems to be a good explanation for eruption ol the flare-associated: CAIEs studied here. but a different mechanism should be considered for EP-associated CAIEs.," In the cases analysed here, we observe that only the flare-associated CMEs undergo residual acceleration, which indicates that flux injection seems to be a good explanation for eruption of the flare-associated CMEs studied here, but a different mechanism should be considered for EP-associated CMEs."
+ Of the three CMESs associated with prominences. the 2010 April 13 and 2010 August 1 were associated with large quiescent polar crown prominences. while the one on 2008 April Ü was associated wilh an active-region prominence.," Of the three CMEs associated with prominences, the 2010 April 13 and 2010 August 1 were associated with large quiescent polar crown prominences, while the one on 2008 April 9 was associated with an active-region prominence."
+ We find that the prominences on 2008 April 9 and 2010 April 13 showed a strong positive acceleration in the CORI FOV. when their heights were close to almost ΕΠ...," We find that the prominences on 2008 April 9 and 2010 April 13 showed a strong positive acceleration in the COR1 FOV, when their heights were close to almost $4\Rsun$."
+ During the same time however. acceleration of the CME LE was decreasing.," During the same time however, acceleration of the CME LE was decreasing."
+ This indicates that even at a height of «ΕΙ... forces acting on the CME and the EP cannot be considered to be the same. as suggested by and Mariciéetal.(2004).," This indicates that even at a height of $4\Rsun$, forces acting on the CME and the EP cannot be considered to be the same, as suggested by \citet{Srivastava.etal2000} and \citet{Maricic.etal2004}."
+. Thus. in this study. from the 3D reconstruction of six CMESs and EPs associated with three of them. we have observed some aspects of (heir acceleration. as detailed above. which were nol previously reported.," Thus, in this study, from the 3D reconstruction of six CMEs and EPs associated with three of them, we have observed some aspects of their acceleration, as detailed above, which were not previously reported."
+ We find that the maximum CME acceleration occurs al a height of less than 21... where earlier. this height was believed to be between 2-42...," We find that the maximum CME acceleration occurs at a height of less than $2\Rsun$ , where earlier, this height was believed to be between $2-4\Rsun$."
+ The bimodal acceleration profile was not observed in EP-associated CMESs. but in only those CMES that were not associated with EPs.," The bimodal acceleration profile was not observed in EP-associated CMEs, but in only those CMEs that were not associated with EPs."
+ Two of the three prominences in the study showed a high and rising value of acceleration at a distanceof almost 415. but the, Two of the three prominences in the study showed a high and rising value of acceleration at a distanceof almost $4\Rsun$ but the
+(local) attractors (Bruni1996).,(local) attractors \cite{B}.
+. In. Section 3 we consider the dynamics of the magnetic field. as it collapses with the matter., In Section 3 we consider the dynamics of the magnetic field as it collapses with the matter.
+ The. magnetizecl dynamical svstem is. νοdimensional., The magnetized dynamical system is five-dimensional.
+ Pancakes are still the attractors. with the magnetic field. squeezed in the pancake plane.," Pancakes are still the attractors, with the magnetic field squeezed in the pancake plane."
+ Note that. as the galaxy is formed. tidal forces are generally expected lo hange the orientation. of the ealactic plane relative to the pancake.," Note that, as the galaxy is formed, tidal forces are generally expected to change the orientation of the galactic plane relative to the pancake."
+ Nevertheless. the confinement. of the field in the pancake plane that we find here. is qualitatively consistent with magnetic field. observations in. numerous spiral and disc galaxies.," Nevertheless, the confinement of the field in the pancake plane that we find here, is qualitatively consistent with magnetic field observations in numerous spiral and disc galaxies."
+ We also provide some quantitative results by relating the growth of the field to that of the matter density contrast., We also provide some quantitative results by relating the growth of the field to that of the matter density contrast.
+ When comparing the shearing to the isotropic collapse. we find that anisotropy can lead. to an appreciable increase in the amplification of the magnetic field.," When comparing the shearing to the isotropic collapse, we find that anisotropy can lead to an appreciable increase in the amplification of the magnetic field."
+ Phese analytical results are in qualitative agreement with those of the earlier mentioned. numerical simulations., These analytical results are in qualitative agreement with those of the earlier mentioned numerical simulations.
+ We interpret. them. as an indication. that the magnetic strengths observed in numerous galaxies and galaxy clusters today. could. have resulted. from. seeds considerably weaker than previous estimates., We interpret them as an indication that the magnetic strengths observed in numerous galaxies and galaxy clusters today could have resulted from seeds considerably weaker than previous estimates.
+" The Zeldovich approximation (Zel'dovich|1970) see also (Shandarin&Zeldovieh1989:Padmanabhan1993) for a review ancl (Buchert1996:Alatarrese1996a:Alatar-rese1996b:Ellis&""sagas2002). for an approach similar to ours] arises from a simple ansatz. which extrapolates to the nonlinear regime a well known result of the linear perturbation theory."," The Zel'dovich approximation \cite{Z} [see also \cite{SZ,P} for a review and \cite{buchert,mataa,matab,ET} for an approach similar to ours] arises from a simple ansatz, which extrapolates to the nonlinear regime a well known result of the linear perturbation theory."
+ In the Eulerian [rame (x. /)) which is comoving with the background. expansion. one defines the rescalecl peculiar velocity field à=cdx/de. where e=a(t) is the background. seale factor.," In the Eulerian frame ${\bf x}, t$ ), which is comoving with the background expansion, one defines the rescaled peculiar velocity field ${\bf \tilde{u}}={\rm d} {\bf x}/{\rm d} a$, where $a=a(t)$ is the background scale factor."
+ Phen. to linear order and ignoring decaving modes. &—Vad. where à=242/3155 is the rescaled peculiar gravitational potential and Lf=4fe.," Then, to linear order and ignoring decaying modes, ${\bf \tilde{u}}=-\nabla_{\bf
+x}\tilde{\varphi}$, where $\tilde{\varphi}=2\varphi/3H^2_0a_0^3$ is the rescaled peculiar gravitational potential and $H=\dot a/a$."
+ The Zeldovich approximation uses this linear result in the rescaled nonlinear continuity. Euler aid Poisson equations. which are where à=Opp is the density contrast.," The Zel'dovich approximation uses this linear result in the rescaled nonlinear continuity, Euler and Poisson equations, which are where $\delta=\delta\rho/\rho_{\rm b}$ is the density contrast."
+" Assuming irrotational motion. we define the peculiar expansion and the peculiar. shear by Ó--Vxc0 and στ=dj,106, respectively (0;= 0/007)."," Assuming irrotational motion, we define the peculiar expansion and the peculiar shear by $\tilde{\Theta}=\nabla_{\bf x}\cdot
+\tilde{\bf u}$ and $\tilde{\sigma}_{ij}=\partial_{(j}
+\tilde{u}_{i)} -{1\over3}\tilde{\Theta}\delta_{ij}$ respectively $\partial_i\equiv\partial/\partial x_i$ )."
+ Then. using convective derivatives. Eqs. (1))," Then, using convective derivatives, Eqs. \ref{eq:db}) )"
+" and (2)) give where fj,=OG—(UE3)5;;] is the Newtonian raceless tidal [field and =σσ."," and \ref{eq:ub}) ) give where $E_{ij}=\partial_j\partial_i\tilde{\varphi}
+-(\partial^2\tilde{\varphi}/3)\delta_{ij}$ is the Newtonian traceless tidal field and $2\tilde{\sigma}^2=\tilde{\sigma}_{ij}\tilde{\sigma}_{ij}$."
+ One can then substitute 072. [rom (3)) into (5))., One can then substitute $\partial^2\tilde{\varphi}$ from \ref{poisson}) ) into \ref{eq:ray}) ).
+ In the linear. or in he shear-free case. Eqs. (4))," In the linear, or in the shear-free case, Eqs. \ref{cont}) )"
+ and. (5)) form a local svstenm describing the evolution of the Ελα along its Low lines., and \ref{eq:ray}) ) form a local system describing the evolution of the fluid along its flow lines.
+ In he general case. the presence of the tidal term £7; in (6)) and the lack of a corresponding evolution equation show hat the above svstem is non-local ane not. closed.," In the general case, the presence of the tidal term $E_{ij}$ in \ref{eq:shear}) ) and the lack of a corresponding evolution equation show that the above system is non-local and not closed."
+ Lt also emphasises the fact that in Newtonian gravity. as opposed to general relativity. one cannot consider a purely initial value sroblem.," It also emphasises the fact that in Newtonian gravity, as opposed to general relativity, one cannot consider a purely initial value problem."
+ Instead. as a consequence of action at a distance. one necessarily needs boundary. conditions 1989)..," Instead, as a consequence of action at a distance, one necessarily needs boundary conditions \cite{BG}."
+ The Zeldovich ansatz. implies that the parentheses in (5)) and. (6)) vanish. leacing to a local svstem of ordinary cillerential equations by eliminating the dependence on 9 and £7).," The Zel'dovich ansatz, implies that the parentheses in \ref{eq:ray}) ) and \ref{eq:shear}) ) vanish, leading to a local system of ordinary differential equations by eliminating the dependence on $\delta$ and $E_{ij}$."
+ Since the shear and tide matrices now commute. the above system can be written in the shear-tide eigenframe ancl reduces to three equations. one for Θ and two [or the independent shear eigenvalues σι. 69.," Since the shear and tide matrices now commute, the above system can be written in the shear-tide eigenframe and reduces to three equations, one for $\tilde{\Theta}$ and two for the independent shear eigenvalues $\tilde{\sigma}_1$, $\tilde{\sigma}_2$."
+ Ht is then straightforward to verily that the solutions of the reduced equations that follow from the Zelldovich ansatz are where A; are the three eigenvalues of the initial tidal Ποιά (Alatarrese1996a)., It is then straightforward to verify that the solutions of the reduced equations that follow from the Zel'dovich ansatz are where $\lambda_i$ are the three eigenvalues of the initial tidal field \cite{mataa}.
+. In particular. sel is the solution of the continuity equation (4)). when O is given by (8)).," In particular, $\delta^{\rm zel}$ is the solution of the continuity equation \ref{cont}) ), when $\tilde{\Theta}$ is given by \ref{expZ}) )."
+ On the other hand. if is the density contrast used in the Poisson equation. then from (3)) and (7)) one gets gol575 as à consequence of the approximation (Nusseretal 19901).," On the other hand, if $\delta^{\rm dyn}$ is the density contrast used in the Poisson equation, then from \ref{poisson}) ) and \ref{ansatz}) ) one gets $\delta^{\rm dyn}=-a\tilde{\Theta}^{\rm zel}\not=\delta^{\rm zel}$ as a consequence of the approximation \cite{nusser}."
+. Note that a negative eigenvalue A; corresponds to collapse along the associated: shear eigen-direction., Note that a negative eigenvalue $\lambda_i$ corresponds to collapse along the associated shear eigen-direction.
+ Also. planar pancakes are solutions corresponding το two vanishing eigenvalues.," Also, planar pancakes are solutions corresponding to two vanishing eigenvalues."
+" For example. A,=0As and Ay«0 describes collapse in the third eigen-direction."," For example, $\lambda_1=0=\lambda_2$ and $\lambda_3<0$ describes collapse in the third eigen-direction."
+ As ds well known. this is not only a solution of the simplified ocal dynamics that follows from the Zeldovich ansatz. but also an exact solution of the full svstem (3))-(6)).," As is well known, this is not only a solution of the simplified local dynamics that follows from the Zel'dovich ansatz, but also an exact solution of the full system \ref{poisson}) \ref{eq:shear}) )."
+ In general one expects at [east one of the A; to be negative and smaller han the other two. and on this basis the generic solution should to tend to a pancake.," In general one expects at least one of the $\lambda_i$ to be negative and smaller than the other two, and on this basis the generic solution should to tend to a pancake."
+ The existence of pancake attractors can be confirmed » a dvnamical svstem approach (Bruni 1996).., The existence of pancake attractors can be confirmed by a dynamical system approach \cite{B}. .
+ Dropping he superseript σον we define the new time variable," Dropping the superscript “zel”, we define the new time variable"
+Carrving out the integral over solid angle [ist leads (o an expression for the power per unit frequency in LAE.,Carrying out the integral over solid angle first leads to an expression for the power per unit frequency in LAE.
+ In the case of synchrotron radiation the integral is well known (Ginzburg&Svrovatski1965)... ancl was derived for this specilic purposeby (1959).," In the case of synchrotron radiation the integral is well known \citep{gs65}, and was derived for this specific purposeby \cite{w59}."
+. The integral needed in the case of LAE is different from that for svnchrotron radiation., The integral needed in the case of LAE is different from that for synchrotron radiation.
+ To avoid the mathematical difficulty. discussed in Sec. 4.4..," To avoid the mathematical difficulty discussed in Sec. \ref{dilemma},"
+ we integrate the enussivily (3.1)) over both frequency aud sold augle. and change the variables of integration lo z—z.(12d0) and z: D]c eu)s," we integrate the emissivity \ref{eta2a}) ) over both frequency and sold angle, and change the variables of integration to $z=z_c(1+\Theta)$ and $z_c$: )= )."
+" By writing AP(z,(14-9))=fdzd(2-—2(1+0))Ai'(:) in (3.3)). and performing the &.-integral over the 6 function. one obtains —"," By writing ${\rm Ai}^2\big(z_c(1+\Theta)\big)=\int\rmd z\,\delta\big(z-z_c(1+\Theta)\big)\,{\rm Ai}^2(z)$ in \ref{eta4}) ), and performing the $\xi_c$ -integral over the $\delta$ function, one obtains ."
+" By writing AP(z,(14-9))=fdzd(2-—2(1+0))Ai'(:) in (3.3)). and performing the &.-integral over the 6 function. one obtains —o"," By writing ${\rm Ai}^2\big(z_c(1+\Theta)\big)=\int\rmd z\,\delta\big(z-z_c(1+\Theta)\big)\,{\rm Ai}^2(z)$ in \ref{eta4}) ), and performing the $\xi_c$ -integral over the $\delta$ function, one obtains ."
+by a shallow decay. tend to have a much lower luminosity during the shallow decline. which nmav eive some hint on the origin of the shallow declines and clislavor (he energy injection models.,"by a shallow decay, tend to have a much lower luminosity during the shallow decline, which may give some hint on the origin of the shallow declines and disfavor the energy injection models."
+ We are very grateful to the referee for insightful comments and Z.-G. Dai. D. Zhang.," We are very grateful to the referee for insightful comments and Z.-G. Dai, B. Zhang, E.-W."
+ Liane. S. Covino. and Z.-P. Jin for helpful communications.," Liang, S. Covino, and Z.-P. Jin for helpful communications."
+ L.5. is grateful to N. Mirabal for proolreading the manuscript ancl M. Caprio for helping on the usage of LevelScheme package., L.S. is grateful to N. Mirabal for proofreading the manuscript and M. Caprio for helping on the usage of LevelScheme package.
+ This work made use of data supplied by the Ulx Science Data Centre at the University of Leicester., This work made use of data supplied by the UK Science Data Centre at the University of Leicester.
+ This work was supported by the National Natural Science Foundation ol China (grants 10673034. and 10621303) and the National 973 Project on Fundamental Researches of China (2007CBS15404 and 2009€!D324800)., This work was supported by the National Natural Science Foundation of China (grants 10673034 and 10621303) and the National 973 Project on Fundamental Researches of China (2007CB815404 and 2009CB824800).
+During the past decade there has been considerable experimental developinent in the determination of neutrino masses and mixings [1].,During the past decade there has been considerable experimental development in the determination of neutrino masses and mixings \cite{1}.
+. Recently. T2Ik experiment |2]. has given unambiguous hints of a relatively large 1-3 mixing angle.," Recently, T2K experiment \cite{2} has given unambiguous hints of a relatively large 1-3 mixing angle."
+ In this light. it is natural to look for models which. naturally. accommodate a non-zero value of reactor mixing angle while the atmospheric mixing angle remains near ils maximal value.," In this light, it is natural to look for models which, naturally, accommodate a non-zero value of reactor mixing angle while the atmospheric mixing angle remains near its maximal value."
+ Recently many papers have appeared which reproduce the relatively large value of the reactor mixing angle [3]., Recently many papers have appeared which reproduce the relatively large value of the reactor mixing angle \cite{3}.
+. There are mainly (wo approaches to explain neutrino mixings: 1) Mass independent textures which lead (o mixing matrices independent of the eigenvalues., There are mainly two approaches to explain neutrino mixings: 1) Mass independent textures \cite{4} which lead to mixing matrices independent of the eigenvalues.
+" The most celebrated example of this category is the Gibimaximal (TBA) [5]. scenario which has been derived. from Family svnnmetiries and predicts a vanishing 1-3B müxing angle A,=0. maximal 2-3D mixing anglefo,=w/t and 1-2 mixing angle (4 = sin !(1//3)."," The most celebrated example of this category is the tribimaximal (TBM) \cite{5} scenario which has been derived from family symmetries and predicts a vanishing 1-3 mixing angle $\theta_{13} = 0$, maximal 2-3 mixing angle$\theta_{23} = \pi/4$ and 1-2 mixing angle $\theta_{12}$ = $\sin^{-1}(1/\sqrt{3})$ ."
+ Non-zero 054 can be accommodated in, Non-zero $\theta_{13}$ can be accommodated in
+"estimate of the density parameter: This result. supporting the fact that the density of non-relativistic matter in (he Universe is well bellow the critical one. is consistent with other independent dvnanmical estimates of QO... such as (that by 2.. who used mean relative peculiar velocity measurements for pairs of galaxies aud obtained Qu=0.30Ht,","estimate of the density parameter: This result, supporting the fact that the density of non-relativistic matter in the Universe is well bellow the critical one, is consistent with other independent dynamical estimates of $\Omm$, such as that by \cite{pairs}, who used mean relative peculiar velocity measurements for pairs of galaxies and obtained $\Omm=0.30^{+0.17}_{-0.07}$."
+ Our findings can be verified in at least (wo wavs., Our findings can be verified in at least two ways.
+ The first would be to examine the growth of ihe 2MASS dipole with redshifts of galaxies as their distance estimates. using proper weights.," The first would be to examine the growth of the 2MASS dipole with redshifts of galaxies as their distance estimates, using proper weights."
+ In ihe near future we cannot however hope for a uniform and deep enough sample of spectroscopic redshifts for the catalog. even if the 2MASS Redshift Survev is continued (as it was planned to reach dy.=12.25 mag. see ?)).," In the near future we cannot however hope for a uniform and deep enough sample of spectroscopic redshifts for the catalog, even if the 2MASS Redshift Survey is continued (as it was planned to reach $K_s=12.25$ mag, see \citealt{Huchra}) )."
+ The calculation may be thus feasible only for photo-2zs. if they are available for the whole 2\IASS NSC.," The calculation may be thus feasible only for 's, if they are available for the whole 2MASS XSC."
+ A promising direction towards this goal may be to cross-correlale 221ASS with the data [rom the Wide-field Infrared Survey Explorer (WISE.?) Chat are currently. partially available (72) ancl are expected to be fully released in 2012.," A promising direction towards this goal may be to cross-correlate 2MASS with the data from the Wide-field Infrared Survey Explorer \citep[WISE,][]{WISE} that are currently partially available \citep{WISErelease} and are expected to be fully released in 2012."
+ The other possible verification method could be for instance to examine the behavior of the differential dipole. (lo compare the increments of the growth with theoretical predictions of the ACDAL moclel.," The other possible verification method could be for instance to examine the behavior of the differential dipole, to compare the increments of the growth with theoretical predictions of the $\Lambda$ CDM model."
+ The authors would like to thank Adi Nusser and. Mare. Davis for useful comments concerning an earlier version of this manuscript. as well as to the referee for valuable input.," The authors would like to thank Adi Nusser and Marc Davis for useful comments concerning an earlier version of this manuscript, as well as to the referee for valuable input."
+ This publication makes use of data products from the Two \Heron All Sky Survey. which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/Calilornia Institute ol Technology. funded by the National Aeronauties and Space Administration and the National Science Foundation; the NASA/IPAC Extragalactic Database (NED). which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Administration and of the NASA/IPAC Infrared Science Archive. which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract," This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation; the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration and of the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract"
+around the z~2 mass-size relation using the Guoetal.(2009) measurements.,"around the $z\sim2$ mass-size relation using the \cite{guo09}
+ measurements."
+" This results in a decrease in the scatter by only ~0.03 dex, and does not affect our conclusions."," This results in a decrease in the scatter by only $\sim0.03$ dex, and does not affect our conclusions."
+" We note that, even within the limited redshift range under consideration, differences in redshift play role: the galaxies in the 1.75<z«2.5 subsample are clearlya smaller than the 1.5«z1.75 galaxies."," We note that, even within the limited redshift range under consideration, differences in redshift play a role: the galaxies in the $1.75 < z < 2.5$ subsample are clearly smaller than the $1.5 < z <
+1.75$ galaxies."
+ This may explain some of the disagreement between studies of high-redshift quiescent , This may explain some of the disagreement between studies of high-redshift quiescent galaxies.
+"In the effective radii found by galaxies.Mancinietal.particular,(2010) for largesome high-redshift quiescent galaxies could be due to the fact that they select galaxies with 1.4«z1.75."," In particular, the large effective radii found by \cite{man10} for some high-redshift quiescent galaxies could be due to the fact that they select galaxies with $1.4
+< z < 1.75$."
+" In this context, part of the size evolution between z2 and z20 could be due to the appearance of young, relatively large quiescent galaxies after z~2 (e.g., vanDokkumetal.2008;; Franxetal. 2008;; Saraccoetal. 2009;; vanderWeletal. 2009;; Cassataetal. 2011))."," In this context, part of the size evolution between $z\sim2$ and $z=0$ could be due to the appearance of young, relatively large quiescent galaxies after $z\sim2$ (e.g., \citealt{dok08}; \citealt{fra08}; \citealt{sar09}; \citealt{wel09}; \citealt{cas11}) )."
+" We note that Saraccoetal.(2009) find evidence for a correlation of galaxy compactness with stellar age, such that the most compact high-redshift quiescent galaxies contain older stellarpopulations than quiescent galaxies that lie close to the ζΞ0 mass-size relation."," We note that \cite{sar09} find evidence for a correlation of galaxy compactness with stellar age, such that the most compact high-redshift quiescent galaxies contain older stellarpopulations than quiescent galaxies that lie close to the $z=0$ mass-size relation."
+" We investigate this correlation in Figure 7,, using rest-frame U—Vcolor as a proxy for galaxy "," We investigate this correlation in Figure \ref{fig:compactness}, , using rest-frame $U-V$color as a proxy for galaxy age."
+We define galaxy compactness as the offset between the age.z~2 galaxy sizes and the z=0 mass-size relation of Shenet (2003):: re/Tez=0=re/(2.88x10° M959).," We define galaxy compactness as the offset between the $z\sim2$ galaxy sizes and the $z=0$ mass-size relation of \cite{she03}: : $r_e/r_{e,z=0} = r_e/(2.88 \times 10^{-6} \times
+M^{0.56})$ ."
+ We find no evidence for a correlation between galaxy compactness and, We find no evidence for a correlation between galaxy compactness and
+mapping of backeround starlight. mapping of müllimeter/subiüllimeter dust contiuuuni enission. and mapping of dust absorption against the bright nid-IR backerouud cussion (Berein&Tatalla2007).,"mapping of background starlight, mapping of millimeter/submillimeter dust continuum emission, and mapping of dust absorption against the bright mid-IR background emission \citep{bertaf07}."
+. A πα characteristic of the deusitv profiles. derived with the above techniques is that they require a central flattening., A main characteristic of the density profiles derived with the above techniques is that they require a central flattening.
+ The density. of low-mass cores is almost constant within radi simaller than 25005000 AU with typical central densities of 10? 10° cm? (Mottectal.1998:Ward-Thommpsonet19001.," The density of low-mass cores is almost constant within radii smaller than $2500 - 5000\,$ AU with typical central densities of $^{5}$ – $^{6}$ $^{-3}$ \citep{motetal98,wamoan99}."
+ A popular approach is to describe these cores as truucated isothermal (Donnor-Ebert) sphere (Ebert1955:Bounor 1956).. that often provides a good fit to the data (Bacmannetal.2001:AlvesKaudori 2005).," A popular approach is to describe these cores as truncated isothermal (Bonnor-Ebert) sphere \citep{ebert55,bonnor56}, that often provides a good fit to the data \citep{bacmanetal01,alvesetal01,kandorietal05}."
+.. Bounor-Ebert spheres are equilibrium solutions of sclf-eravitating gas bounded by external pressure., Bonnor-Ebert spheres are equilibrium solutions of self-gravitating gas bounded by external pressure.
+ ILowever. such deusity structure is not unique.," However, such density structure is not unique."
+ Nunerical calculations of the dyuanical evolution of supersonically iubuleut clouds show that transient cores forming at 1ο stagnation poiuts of convergent flows exlübit similar norpholoey (Ballesteros-Paredesetal.2003:EKlesseuetal.2005:Banerjee 2009)...," Numerical calculations of the dynamical evolution of supersonically turbulent clouds show that transient cores forming at the stagnation points of convergent flows exhibit similar morphology \citep{baklva03,klessenetal05,banerjeeetal09}. ."
+" The situation is ess clear when it comes to hieh-uass cores (Beutleretal 2007). because nost highauass cores studied o date show at least some sigu of star formation or irn out to consist of several sub-condeusatious when observed with high enough resolution ίσιο,Beutheretal.2005:Beuther&Wenning 2009)."," The situation is less clear when it comes to high-mass cores \citep{beutheretal07}, because most high-mass cores studied to date show at least some sign of star formation or turn out to consist of several sub-condensations when observed with high enough resolution \cite[e.g.][]{beutetal05,beuthen09}."
+.. Towever. large-scale survers. ee. as conducted in the CweuuseX reeion (Motteetal.2007). iudicate that Lhieh-auass cores are dn many aspects simular to scaled-up versions of low-mass cores.," However, large-scale surveys, e.g. as conducted in the Cygnus-X region \citep{motteetal07}, indicate that high-mass cores are in many aspects similar to scaled-up versions of low-mass cores."
+ We followed these lines of reasoniue and chose an initial density profile with flat immer core and rt? density profile outside a radius of r=O.5 pe.," We followed these lines of reasoning and chose an initial density profile with flat inner core and $r^{-1.5}$ density profile outside a radius of $r=0.5\,$ pc."
+ Although we beein with smooth rather than turbulent initial conditions. bv the time ionizing radiation begiuus to be emitted. gravitational fragmcutation has already produced substantial density perturbatious.," Although we begin with smooth rather than turbulent initial conditions, by the time ionizing radiation begins to be emitted, gravitational fragmentation has already produced substantial density perturbations."
+ The secoud important parameter is the initial rotation of this core. which defines the total amount of augular moment in the svstem aud hence the radial exteut of the protostellar accretion disk.," The second important parameter is the initial rotation of this core, which defines the total amount of angular momentum in the system and hence the radial extent of the protostellar accretion disk."
+ To assess the relevance of fragmentation-mduced. starvation for high-mass star formation we therefore must courpare our choice of 3}=0.05 for the ratio between rotational aud potential energy iu our 1000AL. core both with observational data of molecular cloud cores and with numerical simulations of core formation.," To assess the relevance of fragmentation-induced starvation for high-mass star formation we therefore must compare our choice of $\beta = 0.05$ for the ratio between rotational and potential energy in our $1000\,M_\odot$ core both with observational data of molecular cloud cores and with numerical simulations of core formation."
+ Our choice of values correspond to a specific angular momentum of j=1.27«107 cues |.," Our choice of values correspond to a specific angular momentum of $j = 1.27 \times 10^{23}\,$ $^2$ $^{-1}$ ."
+ The high-auass cloud. cores observed by Pirogovctal.(2005) with amasses up to a few thousaud solar niasses show specific augular momenta j from a few «107 cm?s 1 up to 1075 cues 15 corresponding to J- up to 0.07.," The high-mass cloud cores observed by \citet{pirogovetal03} with masses up to a few thousand solar masses show specific angular momenta $j$ from a few $\times 10^{22}\,$ $^2$ $^{-1}$ up to $10^{23}\,$ $^2$ $^{-1}$, corresponding to $\beta$ -values up to $0.07$."
+ We took this sample as motivation for our choice of parameters., We took this sample as motivation for our choice of parameters.
+ Interesting in this context is also the large-scale rotational motion of molecular gas around a cluster of hyperconipact rregious in C(20.08-0.1 £N. seen by Galvin-Macdid (2009)., Interesting in this context is also the large-scale rotational motion of molecular gas around a cluster of hypercompact regions in G20.08-0.14 N seen by \citet{gavmadetal09}.
+. This flow has a radius of about 0.5 pc. roughly consistent with our simulation.," This flow has a radius of about $0.5\,$ pc, roughly consistent with our simulation."
+ Similar observations were made bv eto(1990). for CGUI0.6-0.1 aud. for WOLA. It ποσο» to be noted. however. that although the specific angular momenta of observed cores exhibit considerable scatter there is a clear treud of decreasing j with decreasing core mass A.," Similar observations were made by \citet{keto90} for G10.6-0.4 and \citet{welchetal87} for W94A. It needs to be noted, however, that although the specific angular momenta of observed cores exhibit considerable scatter there is a clear trend of decreasing $j$ with decreasing core mass $M$."
+ For example. the mean values in the sample discussed by Goodmanetal.(1993) are jocm6&l0?bens8 3 with masses up to several hundred solar masses. while Casellietal.(2002). studied low-mass cores with typically ouly a few solar masses aud find j=T1079 cues +.," For example, the mean values in the sample discussed by \citet{goodmanetal93} are $j \approx 6 \times 10^{21}\,$ $^2$ $^{-1}$ with masses up to several hundred solar masses, while \citet{casellietal02} studied low-mass cores with typically only a few solar masses and find $j \approx 7 \times 10^{20}\,$ $^2$ $^{-1}$."
+ Tu all cases the typical values for the ratio between rotational energy aud poteutial energy is Jm0.05., In all cases the typical values for the ratio between rotational energy and potential energy is $\beta \lesssim 0.05$.
+ This treud coutinues down to stellar scales., This trend continues down to stellar scales.
+ A binary C star with a orbital period of 3 davs has jwlM oaness t. while the spin of a typical T Tauri star is a few <10! cm? 4.," A binary G star with a orbital period of 3 days has $j \approx 10^{19}\,$ $^2$ $^{-1}$, while the spin of a typical T Tauri star is a few $\times 10^{17}\,$ $^2$ $^{-1}$."
+" Ow own Sum rotates ouly with jzmLOM cuss ο, "," Our own Sun rotates only with $j\approx 10^{15}\,$ $^2$ $^{-1}$ ."
+That means. during the process of star formation most of the iuitial aueular moment is removed from the collapsed object (Bodenleiner1995).," That means, during the process of star formation most of the initial angular momentum is removed from the collapsed object \citep{boden95}."
+. If we look for guidance from ποΊσα] simulations of molecular cloud formation aud fragmentation (MacLow&IEKlesseu2001:lessenetal.2009) we see a very simular treud.," If we look for guidance from numerical simulations of molecular cloud formation and fragmentation \citep{maclow04,klessenetal09} we see a very similar trend."
+ Gammueetal.(2003). find a typical mean specific angular moment of j=4«d107 cues |. while Jappsen&Klesseu(2001). find values in the ranee 102 cues tp<107+ ens D for low-iass cores in the simulations by Ixlesseuetal.(1998) aud EKlesseu&Burk-ert(2001.2000). depending ou the evolutionary state of protostellar collapse.," \citet{gammieetal03} find a typical mean specific angular momentum of $j = 4 \times 10^{22}\,$ $^2$ $^{-1}$, while \citet{jappkles04} find values in the range $10^{20}\,$ $^2$ $^{-1} \lesssim j \lesssim10 ^{21}\,$ $^2$ $^{-1}$ for low-mass cores in the simulations by \citet{klessetal98} and \citet{klesburk01,klesburk00} depending on the evolutionary state of protostellar collapse."
+ See also Tilley&Puditz(2007) and Offueretal.(2008)., See also \citet{tillpudr07} and \citet{offetal08}.
+. Altogether. we find that our initial couditious are consistent with simulations of core formation.," Altogether, we find that our initial conditions are consistent with simulations of core formation."
+ The monolithic collapse model rests on the similarity between the shape of the observed core mass cistribution (Motteetal.1998:Johustoue2000.2006:Ladaetal.2008) and the stellar initial mass function. IMFE (Ikroupa2002:Chabrier2003).," The monolithic collapse model rests on the similarity between the shape of the observed core mass distribution \citep{motetal98,johnstone00,johnstone06,ladaetal08} and the stellar initial mass function, IMF \citep{kroupa02,chabrier03}."
+. Tt assumes a one-to-one relation between the distributions with oulv a constant cficicucy factor separating the two functions., It assumes a one-to-one relation between the distributions with only a constant efficiency factor separating the two functions.
+ Each molecular cloud core collapses to form a single star or at most a close binary svstem. with protostellar cedback processes determining the οποιον of the rocess (Matzuor&Mel&ee2000).," Each molecular cloud core collapses to form a single star or at most a close binary system, with protostellar feedback processes determining the efficiency of the process \citep{matzner00}."
+. This model reduces 16 probleii of the origin of the IME to the problem of determining the mass spectrum of bound cores. although strictly speaking the idea that the IAIF is set bv the nass spectrin of cores js independent of auv particular nodel for the origin of that mass spectra.," This model reduces the problem of the origin of the IMF to the problem of determining the mass spectrum of bound cores, although strictly speaking the idea that the IMF is set by the mass spectrum of cores is independent of any particular model for the origin of that mass spectrum."
+ Arguments o explain the core mass distribution ecnerally rely ou je statistical properties of turbulence (Usiessen2001:Ww 109).. which eeuerate structures with a pure powerlaw nass spectrum.," Arguments to explain the core mass distribution generally rely on the statistical properties of turbulence \citep{klessen01,padnor02,hencha08,hencha09}, which generate structures with a pure powerlaw mass spectrum."
+ The thermal Jeans mass in the cloud iendHuposes the flattening and turu-dowu iu the observed mass spectru., The thermal Jeans mass in the cloud thenimposes the flattening and turn-down in the observed mass spectrum.
+ Towever. there are a muuber of caveats.," However, there are a number of caveats."
+ May of the xestellu: cores found m observational surveys appear ο be stable eutities and thus are uulikelv to be in a state of active star formation., Many of the prestellar cores found in observational surveys appear to be stable entities and thus are unlikely to be in a state of active star formation.
+ In addition. the simple interpretation that one core forms on averageone star. and that all cores contain the same uuuber of thermal Jeans masses.leads to a timescale problem," In addition, the simple interpretation that one core forms on averageone star, and that all cores contain the same number of thermal Jeans masses,leads to a timescale problem"
+"find that neither cut nor 0 can be well constrained on their own, but are largely degenerate.","find that neither $\mu_{\rm cut}$ nor $b$ can be well constrained on their own, but are largely degenerate."
+" Choosing [cut in the range [0.04,0.1] or b in the range [0.8,1.5] gives fits which match the overall subhalo halo mass function reasonably well."," Choosing $\mu_{\rm cut}$ in the range $[0.04, 0.1]$ or $b$ in the range $[0.8,1.5]$ gives fits which match the overall subhalo halo mass function reasonably well."
+ To break the degeneracy between µειι and b in our fits we set b=1.2 and fitted eqn. (1)), To break the degeneracy between $\mu_{\rm cut}$ and $b$ in our fits we set $b = 1.2$ and fitted eqn. \ref{eq:fits}) )
+ by varying just µί and µοιι., by varying just $\muone$ and $\mu_{\rm cut}$.
+" The best-fit values for µί and µειι are: By examining the subhalo mass functions of individual halos, we have found that the main scatter is in the normalisation rather than the shape; halos with a higher than average abundance of low-mass subhalos also tend to have a higher than average abundance of high-mass halos and vice versa."," The best-fit values for $\muone$ and $\mu_{\rm
+cut}$ are: By examining the subhalo mass functions of individual halos, we have found that the main scatter is in the normalisation rather than the shape; halos with a higher than average abundance of low-mass subhalos also tend to have a higher than average abundance of high-mass halos and vice versa."
+" At a given Maup/Ms2oo, there is a weak trend in the abundance of subhaloes with host halo mass."," At a given ${\rm M_{sub}/M_{200}}$, there is a weak trend in the abundance of subhaloes with host halo mass."
+" In the region of overlap in Msup/M2o0, the cluster haloes have a mass fraction of substructures that is about 2596 higher than in the galaxy haloes."," In the region of overlap in ${\rm M_{sub}/M_{200}}$, the cluster haloes have a mass fraction of substructures that is about $25\%$ higher than in the galaxy haloes."
+ The mean of the group sample is intermediate between these two., The mean of the group sample is intermediate between these two.
+" The 1596 difference between group and cluster halo identified by Gaoetal.(2004a) is visible at the smaller mass ratios plotted, Msun/M200107."," The $15\%$ difference between group and cluster halo identified by \cite{g04a} is visible at the smaller mass ratios plotted, ${\rm M_{sub}/M_{200}}
+\sim 10^{-4}$ ."
+" The mass dependence of the subhalo abundance with host halo mass in our simulations is much weaker than a theoretical expectation based on semianalytical modelling of the subhalo population (e.g.,vandenBoschetal.2005).."," The mass dependence of the subhalo abundance with host halo mass in our simulations is much weaker than a theoretical expectation based on semianalytical modelling of the subhalo population \citep[e.g.,][]{van05}. ."
+ 'The mass fractions in subhaloes as a function of relative mass for host haloes of different size are shown in Fig. 3.., The mass fractions in subhaloes as a function of relative mass for host haloes of different size are shown in Fig. \ref{fig:massfrac}.
+ 'The curves are very steep at high values reflecting the fact that most of a typical halo's mass in subhaloes is contributed by a relatively small number of very massive subhaloes., The curves are very steep at high values reflecting the fact that most of a typical halo's mass in subhaloes is contributed by a relatively small number of very massive subhaloes.
+" As shown by Springeletal.(2008a) for the Aquariushaloes, these massive subhaloes tend to be in the outer parts of their parent halo."," As shown by \citet{springel08a} for the Aquariushaloes, these massive subhaloes tend to be in the outer parts of their parent halo."
+ The subhalo mass fraction above a given value of Msup/Me200 depends on the mass of the parent halo., The subhalo mass fraction above a given value of ${\rm M_{sub}}/M_{200}$ depends on the mass of the parent halo.
+" Typically, cluster haloes contain about more subhalo mass than galaxy haloes."," Typically, cluster haloes contain about more subhalo mass than galaxy haloes."
+ The Aquarius galactic haloes resolve substructure with relative mass as small as 107., The Aquarius galactic haloes resolve substructure with relative mass as small as $10^{-6}$.
+ Subhaloes above this mass contain about of the total halo mass within R200., Subhaloes above this mass contain about of the total halo mass within $R_{200}$.
+" For clusters, the resolves subhaloes with relative mass just below 10?."," For clusters, the resolves subhaloes with relative mass just below $10^{-5}$."
+ Subhaloes more massive than this already amount to about of the mass within Roo., Subhaloes more massive than this already amount to about of the mass within $R_{200}$.
+" Note, however, that there is a rather large variance from halo to halo in the subhalo mass fraction above a given value of Maus/M2oo, particularly at the high mass end, where much of the subhalo mass is typically contributed by one or two objects."," Note, however, that there is a rather large variance from halo to halo in the subhalo mass fraction above a given value of ${\rm M_{sub}}/M_{200}$, particularly at the high mass end, where much of the subhalo mass is typically contributed by one or two objects."
+ We expect haloes of amass to contain more subhaloes at earlier times because the earlier counterparts are both less concentrated and dynamically younger., We expect haloes of a to contain more subhaloes at earlier times because the earlier counterparts are both less concentrated and dynamically younger.
+ This trend was seen in the simulations of Gaoetal.(2004a)., This trend was seen in the simulations of \cite{g04a}.
+". To investigate the redshift dependence of the subhalo mass function in our simulations, we use a sample of group haloes of mass [1,3]1013 iin theMS-II."," To investigate the redshift dependence of the subhalo mass function in our simulations, we use a sample of group haloes of mass $[1, 3] \times 10^{13}$ in the."
+" We also use the Aquarius galaxy haloes at z=0 and their main progenitors at z=1, the mean mass of these halo is Maoo=1.19x102571 aat z=0 and Mooo=6.73x10! Mgat z—1."," We also use the Aquarius galaxy haloes at $z=0$ and their main progenitors at $z=1$, the mean mass of these halo is $M_{200} = 1.19 \times 10^{12}$ at $z=0$ and $M_{200} =
+6.73 \times 10^{11}$ at $z=1$."
+ The results are displayed in Fig 4.., The results are displayed in Fig \ref{fig:redshift}.
+" In theMS-II sample, there are a total of 219, 204, 169 and 41 haloes at the four redshifts shown, z—0,0.5,1,2."," In the sample, there are a total of $219$, $204$, $169$ and $41$ haloes at the four redshifts shown, $z=0,~0.5,~1,~2$."
+" Clearly, the subhalo mass function evolves with redshift weakly, but systematically, in the expected sense: at high redshift haloes of a given mass contain more subhaloes."," Clearly, the subhalo mass function evolves with redshift weakly, but systematically, in the expected sense: at high redshift haloes of a given mass contain more subhaloes."
+" For example, the subhalo abundance in groups at z=0 is typically 18% lower than at z— 0.5, 2596 lower than at z—1, and 3096 lower than at z—2."," For example, the subhalo abundance in groups at $z=0$ is typically $18\%$ lower than at $z=0.5$ , $25\%$ lower than at $z=1$, and $30\%$ lower than at $z=2$."
+" The subhalo mass function of galaxy haloes evolves slightly more slowly than this, reflecting the earlier formation epoch of galaxy haloes."," The subhalo mass function of galaxy haloes evolves slightly more slowly than this, reflecting the earlier formation epoch of galaxy haloes."
+ In this case the abundance is only higher at z—1 than at z— 0., In this case the abundance is only higher at $z=1$ than at $z=0$ .
+ We note that using eqn. (1)), We note that using eqn. \ref{eq:fits}) )
+" with a and 6 fixed to —0.94 and 1.2, respectively, fits the data shown at each redshift very well."," with $a$ and $b$ fixed to $-0.94$ and 1.2, respectively, fits the data shown at each redshift very well."
+ 'The corresponding parameters are:For thegalaxy haloes the trend persists down to the, The corresponding parameters are:For thegalaxy haloes the trend persists down to the
+Taking the exterior derivative of tliis expression gives or equivalently last equation implies that dp/p is a conserved quautity.,Taking the exterior derivative of this expression gives or equivalently The last equation implies that $\delta\rho/\rho$ is a conserved quantity.
+ Note that in two dimensions. A is a zero-Lorm and (6.39) has a solution: op where e is an infinitesimal constant.," Note that in two dimensions, ${\cal A}$ is a zero-form and always has a solution: where $\epsilon$ is an infinitesimal constant."
+ Since fdp must vanish. A must have zero mean with respect to p.," Since $\int
+\delta\rho$ must vanish, ${\cal A}$ must have zero mean with respect to $\rho$."
+ This can always be arrauged by taking A+.A—«A>. which does not change lie dynamical systen1.," This can always be arranged by taking ${\cal
+A}\rightarrow {\cal A} - <{\cal A}>$, which does not change the dynamical system."
+ The existence o Ais solutiou is a reflection of the fact that there can be Πο chaos with just two degrees of freedom., The existence of this solution is a reflection of the fact that there can be no chaos with just two degrees of freedom.
+ Any uon-trivial invariant. clistribution p is a distributiο over closed orbits on which A is constant. but does not describe the statistics of a system with a given initial coiiditiou.," Any non-trivial invariant distribution $\rho$ is a distribution over closed orbits on which ${\cal A}$ is constant, but does not describe the statistics of a system with a given initial condition."
+ Iu some instances (in any dimension). the invariant probability distribution p used to rever: engineer a dvuauiical system will not «Tiaracterize individual trajectories because of tlie existe of conserved quantities.," In some instances (in any dimension), the invariant probability distribution $\rho$ used to reverse engineer a dynamical system will not characterize individual trajectories because of the existence of conserved quantities."
+ This does uot necessarily exclude chaos. unless the number of couserved quantities is too large.," This does not necessarily exclude chaos, unless the number of conserved quantities is too large."
+ The clistribttion p can then be interpreted as a product of distributi over couserved quantities (which cau ye arbitrary) and a distribution which is depeudeut ou dynamics., The distribution $\rho$ can then be interpreted as a product of distribution over conserved quantities (which can be arbitrary) and a distribution which is dependent on the dynamics.
+ As a word of caution. uot all coiserved quantities associated with non-uniqueness of invariant distribuion p impose strong constraints on trajectories.," As a word of caution, not all conserved quantities associated with non-uniqueness of the invariant distribution $\rho$ impose strong constraints on trajectories."
+ Suppose for example thi dynamical systet1 acdiaits two iuvariaut distributious py and po with different non-overlappi but contiguous domaius of support., Suppose for example that a dynamical system admits two invariant distributions $\rho_1$ and $\rho_2$ with different non-overlapping but contiguous domains of support.
+ Tien there is a class of nou-ergodic distributious of the foZU ipit(1—o)ps wlere QSarx I.," Then there is a class of non-ergodic distributions of the form $x\rho_1 +
+(1-x)\rho_2$ where $0\le x\le 1$ ."
+" The associated conserved quantity is (pj—pa)/(p,+pa). which is 1 in the region of support of py and —1 in the region of support of po."," The associated conserved quantity is $(\rho_1-\rho_2)/(\rho_1+\rho_2)$, which is $1$ in the region of support of $\rho_1$ and $-1$ in the region of support of $\rho_2$."
+ The only. coustraiut this conserved quantity places ou trajectories is that they can not go from the interior of region ] to the interior «X region 2. or vice versa.," The only constraint this conserved quantity places on trajectories is that they can not go from the interior of region 1 to the interior of region 2, or vice versa."
+ C'ouskler the static Fokker-Planck equation for a dynamical system in tlie preseuce of a Craussiau raucdoin force: One cau always (locally)write, Consider the static Fokker-Planck equation for a dynamical system in the presence of a Gaussian random force; One can always (locally)write
+nueht sugeest that these componcuts share a comuuou history (or origin).,might suggest that these components share a common history (or origin).
+ In the next section we describe some scenarios that nüeht have happened when darf galaxies interacted with the ceutral cluster galaxy., In the next section we describe some scenarios that might have happened when dwarf galaxies interacted with the central cluster galaxy.
+ What are the possible consequences. when dwarf galaxies of differcut types interact with the central. ealaxy. especially with respect to the formation of a cD halo aud a rich GCS?," What are the possible consequences, when dwarf galaxies of different types interact with the central galaxy, especially with respect to the formation of a cD halo and a rich GCS?"
+ We make a clistinction between two nuin (1) the iufall of eas-poor chwarfs. for example dw ellipticals. where oulv the existing stellar component is involved iu the interaction process. and (2) the iufall of eas-rich dwarfs. where the interaction of the gasons ji fo be considered and aight plav an Huportant role iu the formation of new stellar populations A further sub-division of these cases (a) the dwarf galaxy will be totally dissolved in the interaction (b) only parts of the dwarf galaxy (for example gas and/or oelobular clusters) will be stripped during the passage through the ceutral cluster (c) the dwarf galaxv neither loses gas nor stars clusters to the cluster center. but might chanee its morphological shape because of tidal interactions (for exaniple getting more conipact or splitting into two).," We make a distinction between two main (1) the infall of gas-poor dwarfs, for example dwarf ellipticals, where only the existing stellar component is involved in the interaction process, and (2) the infall of gas-rich dwarfs, where the interaction of the gas has to be considered and might play an important role in the formation of new stellar populations A further sub-division of these cases (a) the dwarf galaxy will be totally dissolved in the interaction (b) only parts of the dwarf galaxy (for example gas and/or globular clusters) will be stripped during the passage through the central cluster (c) the dwarf galaxy neither loses gas nor stars nor clusters to the cluster center, but might change its morphological shape because of tidal interactions (for example getting more compact or splitting into two)."
+ Iu the next subsectious woe discuss the possible consequences for the different. cases. (, In the next subsections we discuss the possible consequences for the different cases. (
+1a): in this case the stellar population of he dwarf ealaxy will be disrupted in tidal tails aud the stellar ight will be sincared out in the otential well of the cluster ceuter.,1a): in this case the stellar population of the dwarf galaxy will be disrupted in tidal tails and the stellar light will be smeared out in the potential well of the cluster center.
+ Most affected. by this process are the faintest dEs (or dSplis. Thompson Gregory 1993)).," Most affected by this process are the faintest dEs (or dSphs, Thompson Gregory \cite{thom}) )."
+ Ta clusters with a low velocity dispersion or at the bottoms ofa loca potcutial wel in avich cluster (Zabludoff et al. 199001).," In clusters with a low velocity dispersion or at the bottom of a local potential well in a rich cluster (Zabludoff et al. \cite{zabl}) ),"
+ the light of several dissolved. cdiwarts way form an extended. diffuse cD halo.," the light of several dissolved dwarfs may form an extended, diffuse cD halo."
+ Existing GCs of the dwarfs will survive aud contribute to the central CCS., Existing GCs of the dwarfs will survive and contribute to the central GCS.
+ Iu the Local Croup. an exiuuple for this scenario may be the Sagittarius dSpli which is dissolving iuto our Calaxy adding 1 new GCs to the GCS of the Milkv. War (Da Costa Avmandroff 19051).," In the Local Group, an example for this scenario may be the Sagittarius dSph which is dissolving into our Galaxy, adding 4 new GCs to the GCS of the Milky Way (Da Costa Armandroff \cite{daco95}) )."
+ ILowever. only few dwart galaxies with a very rich GCS compared to their Iuuinositv are kuown (Miller et al. 1998.. ," However, only few dwarf galaxies with a very rich GCS compared to their luminosity are known (Miller et al. \cite{mill}, ,"
+Durrell et al. 1996))., Durrell et al. \cite{durr}) ).
+ Tn Sect., In Sect.
+ 6 we estimate under which couditious the accretion of gas-poor dwarfs auc their GCS can increase Sy of a central GCS., 6 we estimate under which conditions the accretion of gas-poor dwarfs and their GCS can increase $S_N$ of a central GCS.
+ Finally. the uuclei of dE.Ns can survive the dissolution of their parcut ealaxy and iav appear as GCs (Ziunecker ot al. 1988..," Finally, the nuclei of dE,Ns can survive the dissolution of their parent galaxy and may appear as GCs (Zinnecker et al. \cite{zinn},"
+ Dassino ct al 1991))., Bassino et al \cite{bass}) ).
+ The nuclear magnitudes of all Virgo dE.Ns (Bineecl Cameron 1991)). for example. fall in the magnitude surface brightucss sequence is defined by the GCs (c.g. BingeeliOO 1991)). (1b): like iu the case la the stripped stars aud GCs will be distributed around the ceutral galaxy.," The nuclear magnitudes of all Virgo dE,Ns (Binggeli Cameron \cite{bing91}) ), for example, fall in the magnitude – surface brightness sequence is defined by the GCs (e.g. Binggeli \cite{bing94}) (1b): like in the case 1a the stripped stars and GCs will be distributed around the central galaxy."
+ Iu this case the question arises ou how large the nuuber of stripped GCs is compared to the Iuninositv of the stripped stellar liebt., In this case the question arises on how large the number of stripped GCs is compared to the luminosity of the stripped stellar light.
+ If GCs could be stripped from regions with a high local Soy. this would also increase Soy of he ceutral GCS.," If GCs could be stripped from regions with a high local $S_N$, this would also increase $S_N$ of the central GCS."
+ According to model calculations by. Muzzio ct al. (198L..," According to model calculations by Muzzio et al. \cite{muzz84},"
+ see also review by Muzzio 1987)). the tidal accretion of CCS and stars can be an important process m the dynamical evolution of GCSs in galaxy clusters.," see also review by Muzzio \cite{muzz87b}) ), the tidal accretion of GCs and stars can be an important process in the dynamical evolution of GCSs in galaxy clusters."
+ Iu some galaxies the CCS is more extended than the nuderlving stellar light. which has the consequence that 1ο local Sy increases with ealactocentric distance: for sxaunple NGC 1172 has a elobal Sy of 5.5 and a local Sy arecr than 30 at 90 kpe (MeLbaughllin et al. 199 1)).," In some galaxies the GCS is more extended than the underlying stellar light, which has the consequence that the local $S_N$ increases with galactocentric distance; for example NGC 4472 has a global $S_N$ of 5.5 and a local $S_N$ larger than 30 at 90 kpc (McLaughlin et al. \cite{mcla94b}) )."
+ Forbes al. (1997)), Forbes et al. \cite{forb97}) )
+ and Ivissler-Patig et al. (1999)), and Kissler-Patig et al. \cite{kiss99a}) )
+ sueeest that je stripping of the outermios GCs and stars from such vealaxy uaturally increases the Sy of the ceutral GCS., suggest that the stripping of the outermost GCs and stars from such a galaxy naturally increases the $S_N$ of the central GCS.
+ It would be interesting to investigate whether this is also rue for the GC'Ss of dwarf galaxies., It would be interesting to investigate whether this is also true for the GCSs of dwarf galaxies.
+ Furthermore. it would be imteresting to know how the idal stripping process changes the shape of the remaining ealaxv.," Furthermore, it would be interesting to know how the tidal stripping process changes the shape of the remaining galaxy."
+ Kroupa (1997)) simulated the interaction of a spherical low-nass ealaxy with a massive ealactic halo and found that the model remnants share the properties of dwarf spheroidals., Kroupa \cite{krou}) ) simulated the interaction of a spherical low-mass galaxy with a massive galactic halo and found that the model remnants share the properties of dwarf spheroidals.
+ On the other haud. N32 may be au exinuple for a tidally compressed renmnant. whose CC's have heen stripped (c.g. Faber 1973... Cepa Beclanan 1985)). (lc) in this case the dwart galaxy does not contribute to the formation of cD halo aud ceutzal CCS.," On the other hand, M32 may be an example for a tidally compressed remnant, whose GCs have been stripped (e.g. Faber \cite{fabe73}, Cepa Beckman \cite{cepa}) (1c): in this case the dwarf galaxy does not contribute to the formation of cD halo and central GCS."
+ Iowever. as in 1b one nüght speculate about the chauge of the morphological shape after a passage of the galaxy through the cluster Note that. except in their unclei. the metallicity of GCs in dEs as well as the metallicity of the bulk of their stars is very low 2.5<|Fe/H|—1.0 dex. see the review on Local Group dwarts by Hodge 199 1).," However, as in 1b one might speculate about the change of the morphological shape after a passage of the galaxy through the cluster Note that, except in their nuclei, the metallicity of GCs in dEs as well as the metallicity of the bulk of their stars is very low $-2.5 < [Fe/H] < -1.0$ dex, see the review on Local Group dwarfs by Hodge \cite{hodg94b}) )."
+" Therefore. stripped GCs from these galaxies will only coutribute to the moetal- populationof the ceutral GCS,"," Therefore, stripped GCs from these galaxies will only contribute to the metal-poor populationof the central GCS."
+ Aud accordingly. the cD halo should have quite a blue color.," And accordingly, the cD halo should have quite a blue color."
+Since there is a direct. traceoll betweet extra constaut [lux from the accretion disk (or any other source) aud the orbital inclination inferred [rom the amplitude of the ellipsoidal variations. we ran a tinal set of models to determiue iow the inclination changes as disk [lux is adde to the light eurve.,"Since there is a direct tradeoff between extra constant flux from the accretion disk (or any other source) and the orbital inclination inferred from the amplitude of the ellipsoidal variations, we ran a final set of models to determine how the inclination changes as disk flux is added to the light curve."
+ To do this. we fit a Ix star uodel (Ts = 1100 Is. 3 = 0.08) to the 1996 Jautary data ancl adde increasing amounts of disk lux. specified in the Iight curve syuthesis progranLas a fraction of the tota licht curve [lux at o 0.25.," To do this, we fit a K star model $_{2}$ = 4100 K, $\beta$ = 0.08) to the 1996 January data and added increasing amounts of disk flux, specified in the light curve synthesis program as a fraction of the total light curve flux at $\phi$ = 0.25."
+ Table 5 shows the derived orbital inclination as the clisk conribution increases., Table \ref{tab_third} shows the derived orbital inclination as the disk contribution increases.
+ For no clisk contribution. we recover the 38° inclination foμιά in Section 3.1..," For no disk contribution, we recover the $\arcdeg$ inclination found in Section \ref{sec_star}. ."
+ 'To alow orbital inclinatiois approaching τοῦ — te upper limit set by the lack of eclipses — the disk would ueed to be contribute of the H-baud flux., To allow orbital inclinations approaching $\arcdeg$ – the upper limit set by the lack of eclipses – the disk would need to be contribute of the H-band flux.
+ Table 5 also shows that if he accretion disk contribution it the H-baud is low. large fuctuaious in the disk flux are needed o explain the chauge in the amXitude of the © = 0.25 peak cisCussed in Section 3.+)1l: the 6-7* difference in the inclilations of our secondary star model fits to he three H light curves implies luctuatious of in the coi(αμαig flux at the o = 0.25 »eak.," Table \ref{tab_third} also shows that if the accretion disk contribution in the H-band is low, large fluctuations in the disk flux are needed to explain the change in the amplitude of the $\phi$ = 0.25 peak discussed in Section \ref{sec_star}: the $\arcdeg$ difference in the inclinations of our secondary star model fits to the three H light curves implies fluctuations of in the contaminating flux at the $\phi$ = 0.25 peak."
+ Based ou our fits to the H-b:ALC --ielit curves. we cau cousrain the inclination i A0620-00 to BAep«4<τοῦ.," Based on our fits to the H-band light curves, we can constrain the inclination in A0620–00 to $38\arcdeg \leq i \leq 75\arcdeg$."
+" The upper limit oi ile inclination is quite str Wa. light curve models with 7275° show both primary auc secoidary lipse features no ""eel u the data."," The upper limit on the inclination is quite strict, as light curve models with $i \geq 75\arcdeg$ show both primary and secondary eclipse features not seen in the data."
+ The lowe‘limit on the inclination is also a fairly strict Hi., The lower limit on the inclination is also a fairly strict limit.
+ The amplitude «X the --elit ονο noculation was larger in 1996 December aud 1995 Dec‘elu lan in 1996 January. iudicatiig that t DN ower limit on the inclination derived from the alter is a probable uncerestimate of the true binary inclination caused by dilution of the ellipsoical inocdulatiou by au acereion disς [lux component.," The amplitude of the light curve modulation was larger in 1996 December and 1995 December than in 1996 January, indicating that the $\arcdeg$ lower limit on the inclination derived from the latter is a probable underestimate of the true binary inclination caused by dilution of the ellipsoidal modulation by an accretion disk flux component."
+ From «etermiuations of the orbital period. mass ratio and r:nial velocity senmi-ainXitude of the Is star. Ma‘sh.Robinson&Wood(1991). derived a mass M4=(3.09+snooE Που the cor1yaet star in 0620-00.," From determinations of the orbital period, mass ratio and radial velocity semi-amplitude of the K star, \cite{marsh1994} derived a mass $M_{1} = (3.09\pm0.09)\ \sin^{-3}i$ for the compact star in A0620-00."
+ Combined with our limits on he inclination. this imits the mass of the cupact star to lie in the rauge 3.)<AL13.6M.," Combined with our limits on the inclination, this limits the mass of the compact star to lie in the range $3.3 \leq M_{1} \leq
+13.6$."
+.. The l.»vwer limit is close to. but remains arger than. the maxim mass of a uuiformly rotating neutror star with the stillest equation of state. ~3.2 citeplriecinar 1980.. and it is well above the παν mass of a nonrotatiug neutron star. 1.5 — 2.5 Citepakimall99e..," The lower limit is close to, but remains larger than, the maximum mass of a uniformly rotating neutron star with the stiffest equation of state, $\sim$ 3.2 \\citep{friedman1986}, and it is well above the maximum mass of a nonrotating neutron star, 1.8 – 2.5 \\citep{akmal1998}."
+ Haswelletal.(1993) identified a sharp dip in their UBVR light curves of A0690-00 as au eclipse of the Ix star by a large accretion disk (at or near the maximum radius lora circular disk insice the Roche lobe around the compact star)., \citet{haswell1993} identified a sharp dip in their UBVR light curves of A0620-00 as an eclipse of the K star by a large accretion disk (at or near the maximum radius fora circular disk inside the Roche lobe around the compact star).
+ I£ this identification is correct. the inclination of," If this identification is correct, the inclination of"
+"NIE molecular ""M",NE molecular cloud.
+ ‘Towards the Triple Core region theNanton? 72 n21) emission is resolved. into two peaks assocated. with IMiirregions €C5.89-3.80X and D respectively., Towards the Triple Core region the Nanten2 $^{12}$ (2–1) emission is resolved into two peaks assocated with the regions G5.89-3.89A and B respectively.
+ Core 5 NW. and. SW. are. also traced. by CO (21) enhancements. which are also clearly seen. in CO (QI0) (sim&Koo2003).," Core 5 NW and SW are also traced by $^{12}$ (2–1) enhancements, which are also clearly seen in $^{13}$ (1–0) \citep{kimkoo2003}."
+ Our Mopra emission also appears to resolve peaks associated with NI;5502SOA and. D. and Cores 5 NE and SW.," Our Mopra $_3$ emission also appears to resolve peaks associated with G5.89-3.89A and B, and Cores 5 NE and SW."
+" Figure 7 presents the infrared. (110) GLIAIDSE anc MIPSCAL three-colour μαι). image of the W28 region with Alopra NIL;""(( hankand LL. TeV eanuna-ray contours."," Figure \ref{fig:IRdata} presents the infrared (IR) GLIMPSE and MIPSGAL three-colour $\mu$ m) image of the W28 region with Mopra $_{3}$ (1,1) and H.E.S.S. TeV gamma-ray contours."
+ These LR are tracers of polvevclic aromatic hydrocarbons (PALIs) ancl dust emission. revealing the complexity of the W2s region in hosting several star formation ancl regions in quite likely several different evolutionary stages.," These IR bands are tracers of polycyclic aromatic hydrocarbons (PAHs) and dust emission, revealing the complexity of the W28 region in hosting several star formation and regions in quite likely several different evolutionary stages."
+ Table 3. sumamarises the likely counterparts identifed with our detected cores., Table \ref{tab:coreID} summarises the likely counterparts identifed with our detected cores.
+ Further discussion. of the. individual cores and comparison of previous molecular line studies toward them can be found later in re[sec:coredisceussion.., Further discussion of the individual cores and comparison of previous molecular line studies toward them can be found later in \\ref{sec:corediscussion}.
+ We describe here procedures. used. to. estimate the gas parameters such as optical depth. temperature. mass. density. ancl related dynamical. information towards the NIL; cores.," We describe here procedures used to estimate the gas parameters such as optical depth, temperature, mass, density, and related dynamical information towards the $_3$ cores."
+ Vablest (with statistical errors in the online appendix) and 5 summarise these results utilising the NIE; CC.1) and (2.2) spectra for an analysis assuming various core sizes. and additionall results treating Core 2. Triple Core and Core 5 as extended. clouds well bevond the size of the 2' PWILM beam.," \ref{tab:gas_parameters} (with statistical errors in the online appendix) and \ref{tab:mass_dens} summarise these results utilising the $_3$ (1,1) and (2,2) spectra for an analysis assuming various core sizes, and additional results treating Core 2, Triple Core and Core 5 as extended clouds well beyond the size of the $^\prime$ FWHM beam."
+ We also apply a detailed radiative transfer model to Core 2. given its apparent high gas temperature ancl non-thermal energy.," We also apply a detailed radiative transfer model to Core 2, given its apparent high gas temperature and non-thermal energy."
+ The ENIM. of the NIL; main line. ενΌροι is α uselul measure of the total energv associated with the core or clump.," The FWHM of the $_{3}$ main line, $\Delta\,v_{1/2}$, is a useful measure of the total energy associated with the core or clump."
+ Droader lines will result. from regions with higher temperatures or some additional dynamics., Broader lines will result from regions with higher temperatures or some additional dynamics.
+" “Phe purely Maxwell-Boltzmann thermal lino width FWIAL Aei expected from. a gas at temperature Z' is given by where & is Boltzmann's constant and mg, is the mass of the NIL; molecule."," The purely Maxwell-Boltzmann thermal line width FWHM $\Delta\,v_{\rm th}$ expected from a gas at temperature $T$ is given by where $k$ is Boltzmann's constant and $m_{\tiny{{\rm NH}}_{3}}$ is the mass of the $_3$ molecule."
+ For example. a thermal line width of ~O.16 is obtained for a temperature of LOW. as might be expected. in typically cold dense NIL; cores (Ho1983).," For example, a thermal line width of $\sim$ is obtained for a temperature of K, as might be expected in typically cold dense $_{3}$ cores \citep{hotownes1983}."
+.- The line ENLEM. Areyeo. of each core was estimated from a Gaussian fit to the central or main peak of the emission. with aclelitional Gaussians to fit each of the four satellite lines. which are generallyresolved. in the (1.1) transition.," The line FWHM, $\Delta\,v_{1/2}$, of each core was estimated from a Gaussian fit to the central or main peak of the emission with additional Gaussians to fit each of the four satellite lines, which are generallyresolved in the (1,1) transition."
+ This five-Gaussian [it function can be seen applied to the Core 1: and Core 2 spectra in Figure &.., This five-Gaussian fit function can be seen applied to the Core 1 and Core 2 spectra in Figure \ref{fig:sample_spectra}.
+ Vhe results in Table 4. show that Nepes for all cores is considerably. wider than that expected. from purely thermal broadening. suggesting additional non-thermal or kinetic energy. which dominates over broadening from the instrumental response. which is considered. negligible.," The results in Table \ref{tab:gas_parameters} show that $\Delta\,v_{1/2}$ for all cores is considerably wider than that expected from purely thermal broadening, suggesting additional non-thermal or kinetic energy which dominates over broadening from the instrumental response, which is considered negligible."
+ Vhe (1.1) satellite lines were clearly resolved in. all cores except in Core 2 which is instrinsically very. broad. leading to blending of the main ancl satellite lines (see Figure S)).," The (1,1) satellite lines were clearly resolved in all cores except in Core 2 which is instrinsically very broad, leading to blending of the main and satellite lines (see Figure \ref{fig:sample_spectra}) )."
+ Using the relative brightness temperatures of the main ancl satellite peaks. the optical depth can be derived. for each (J. AN) inversion transition by numerically solving Panation 2 of Barrettetal.(1977).," Using the relative brightness temperatures of the main and satellite peaks, the optical depth can be derived for each $J,K$ ) inversion transition by numerically solving Equation 2 of \citet{barrett}."
+. A weighted: average (J.Iv1) (with weight given as l/peak-error according to the Gaussian. fit) of the four satellite-derived optical depths is then calculated. ancl the total optical depth τιν of the transition is estimated by accounting for the fraction of intensity in the main line (where. f(1.1)20.502.. and [(2.2320.79ti).," A weighted average $\bar{\tau}(J,K,m)$ (with weight given as 1/peak-error according to the Gaussian fit) of the four satellite-derived optical depths is then calculated, and the total optical depth $\tau_{\rm J,K}^{\rm tot}$ of the transition is estimated by accounting for the fraction of intensity in the main line (where $f$ (1,1)=0.502, and $f$ (2,2)=0.796)."
+ Without needing to resolve the (2.2) satellite lines. the main line optical depth of the (1.1) transition can be used to infer the main line optical depth. of the (2.2) transition using the method of Ungerechtsctal.(1986," Without needing to resolve the (2,2) satellite lines, the main line optical depth of the (1,1) transition can be used to infer the main line optical depth of the (2,2) transition using the method of \citet{unger}."
+) As [or the (1.1) transition. a weighted average may be used to derive the total (2.2) optical depth τος.," As for the (1,1) transition, a weighted average may be used to derive the total (2,2) optical depth $\tau_{\rm 2,2}^{\rm tot}$."
+ Phe line EAWHIAL is of eritical consideration. when determining the emperature of the gas. as a key assumption in this step is that both the (1.1) and (2.2) emission. probe the same volume of gas.," The line FWHM is of critical consideration when determining the temperature of the gas, as a key assumption in this step is that both the (1,1) and (2,2) emission probe the same volume of gas."
+ The similarity of FWIIM for the (1.1) and (2.2) transitions in Table 4. does however suggest that xh transitions originate from the same general volume of eas. allowing the rotational temperature to be caleulated using the method. of Ungerechtsctal.(1986)...," The similarity of FWHM for the (1,1) and (2,2) transitions in Table \ref{tab:gas_parameters} does however suggest that both transitions originate from the same general volume of gas, allowing the rotational temperature to be calculated using the method of \citet{unger}. ."
+ Generally he rotational temperature is an uncerestimate of the kinetic, Generally the rotational temperature is an underestimate of the kinetic
+with noise-excited features to produce an overall response which is lager than the time-asvmptotic predictions described above.,with noise-excited features to produce an overall response which is larger than the time-asymptotic predictions described above.
+ Because the modes are already excited by noise. they may be entrained and strenethec by the external perturbation.," Because the modes are already excited by noise, they may be entrained and strengthed by the external perturbation."
+" ""This is similar to other noise amplification mechanisms found in Nature. e.g. mechanism observed. in neurobiologv (e.g. Collins. Chow Imboll 1995))."," This is similar to other noise amplification mechanisms found in Nature, e.g. mechanism observed in neurobiology (e.g. Collins, Chow Imhoff \nocite{CoCI:95}) )."
+ Moreover. there are many sources of inhomogeneity in real galaxies such as star clusters. gas clouds. unmixed streams from clwarl dissolution. and other sources of perturbations such as tidal distortions from interactions with a host galaxy. cluster and near-neighbor interactions.," Moreover, there are many sources of inhomogeneity in real galaxies such as star clusters, gas clouds, unmixed streams from dwarf dissolution, and other sources of perturbations such as tidal distortions from interactions with a host galaxy cluster and near-neighbor interactions."
+ In fact. the noise in the simulation. presented here corresponds to that. produced. by a halo of 10AL. black holes (e.g. Lacey Ostriker 1985)).," In fact, the noise in the simulation presented here corresponds to that produced by a halo of $10^6M_\odot$ black holes (e.g. Lacey Ostriker \nocite{LaOs:85}) )."
+ Altogether. it seems Likely that observable elfects on. disks due to halo interactions will be stronger than the predictions in relsee:results..," Altogether, it seems likely that observable effects on disks due to halo interactions will be stronger than the predictions in \\ref{sec:results}."
+ This paper summarizes the οσοι of a satellite companion on a galaxian disk ορούσα in a responsive (or five) halo., This paper summarizes the effect of a satellite companion on a galaxian disk embedded in a responsive (or ) halo.
+ Because the mechanisms described. here involve. multiple time and length scales. the interaction 1s complex and depends on the details of each dynamical system: satellite orbit. halo and disk structure.," Because the mechanisms described here involve multiple time and length scales, the interaction is complex and depends on the details of each dynamical system: satellite orbit, halo and disk structure."
+ The general conclusions ancl expectations are enumerated. below: The halo wake (c.g. Figs., The general conclusions and expectations are enumerated below: The halo wake (e.g. Figs.
+ 2. and 3)) ῥίανς a critical role in producing a warp and is an observable consequence of the massive dark. halo., \ref{fig:halowake1} and \ref{fig:halowake2}) ) plays a critical role in producing a warp and is an observable consequence of the massive dark halo.
+ A warp. survey combined with recent satellite surveys (Zaritsky et al. 1993)), A warp survey combined with recent satellite surveys (Zaritsky et al. \nocite{ZRFW:93}) )
+ may provide additional statistical evidence for the massive halo hypothesis., may provide additional statistical evidence for the massive halo hypothesis.
+ Several possibly important interactions have been ignored in this study and present topics for further research., Several possibly important interactions have been ignored in this study and present topics for further research.
+ First. the disk is model is infinitely thin: including the three-dimensional structure will damp the warp producing modes but this damping is likely to be small at large scales (Weinberg. 1991)).," First, the disk is model is infinitely thin; including the three-dimensional structure will damp the warp producing modes but this damping is likely to be small at large scales (Weinberg \nocite{Wein:91b}) )."
+ Second. the disk respond. to the halo wake and the halo responds to the two-dimensional. disk distortion but not the three-cdimensional one.," Second, the disk respond to the halo wake and the halo responds to the two-dimensional disk distortion but not the three-dimensional one."
+ Therefore the dynamical friction. against the halo explored. by Nelson ‘Tremaine (1995)) is not included., Therefore the dynamical friction against the halo explored by Nelson Tremaine \nocite{NeTr:95}) ) is not included.
+ The caleulational method used. here is tractable because of the assumption that all transients have mixed. away., The calculational method used here is tractable because of the assumption that all transients have mixed away.
+ As they mix. they produce Iluctuations on many. scales.," As they mix, they produce fluctuations on many scales."
+ Latrinsie halo inhomogencities such as clouds. clusters and massive black holes are also a source of noise.," Intrinsic halo inhomogeneities such as clouds, clusters and massive black holes are also a source of noise."
+ Together. this noise may seed: interesting observable features. such as inner bars and arms. which are not part of the long-term wake.," Together, this noise may seed interesting observable features, such as inner bars and arms, which are not part of the long-term wake."
+ In addition. as described in," In addition, as described in"
+Therefore the present results extend beyond the stellar wind case and may be applicable to both long and short GRBs.,Therefore the present results extend beyond the stellar wind case and may be applicable to both long and short GRBs.
+ Interestingly. recent Swift observations have revealed a rather complex early X-ray afterglow light curve in a subset of GRBs. with a fast decay at @..~LOOs followed by a form of plateau that joins a more standard light curve at later times 10! s (Nousek et al.," Interestingly, recent Swift observations have revealed a rather complex early X-ray afterglow light curve in a subset of GRBs, with a fast decay at $t_{\rm obs}\sim100\,$ s followed by a form of plateau that joins a more standard light curve at later times $\gtrsim
+10^4\,$ s (Nousek et al."
+ 2006. O'Brien et al.," 2006, O'Brien et al."
+ 2006)., 2006).
+ High latitude emission is considered as a possible explanation for the steep decay phase. although modelling the plateau phase with the afterglow brings in additional constraints on the overall GRB model (e.g. Panaitescu 2007).," High latitude emission is considered as a possible explanation for the steep decay phase, although modelling the plateau phase with the afterglow brings in additional constraints on the overall GRB model (e.g. Panaitescu 2007)."
+ The present scenario could account for two of these observed features — the initial fast decay and the late time recovery — but it does not explain the emergence of the plateau., The present scenario could account for two of these observed features – the initial fast decay and the late time recovery – but it does not explain the emergence of the plateau.
+ The following briefly addresses these issues in turn., The following briefly addresses these issues in turn.
+ Regarding the fast decay phase. the present scenario predicts an exponential decay and a clear spectral transition from hard to soft as the peak of the emission exits the X-ray band.," Regarding the fast decay phase, the present scenario predicts an exponential decay and a clear spectral transition from hard to soft as the peak of the emission exits the X-ray band."
+ One does not therefore expect a perfectly smooth transition from the prompt emission to the fast decay phase., One does not therefore expect a perfectly smooth transition from the prompt emission to the fast decay phase.
+ As discussed in Sec. ??..," As discussed in Sec. \ref{sec:lcurve},"
+ additional theoretical developments are required to provide a detailed light curve around 1008s. Nevertheless. it is of interest to note that Sakamoto et al. (," additional theoretical developments are required to provide a detailed light curve around $100\,$ s. Nevertheless, it is of interest to note that Sakamoto et al. ("
+2007) have reported evidence for an exponential decay component on top of a powerlaw decaying component in the early X-ray light curve.,2007) have reported evidence for an exponential decay component on top of a powerlaw decaying component in the early X-ray light curve.
+ Furthermore Zhang et al. (, Furthermore Zhang et al. (
+2007) have observed a pronounced hard-to-soft spectral transition during the fast decay in two thirds of GRBs that show a fast decay (see also Yonetoku et al.,2007) have observed a pronounced hard-to-soft spectral transition during the fast decay in two thirds of GRBs that show a fast decay (see also Yonetoku et al.
+ 2008)., 2008).
+ Their phenomenological model involves an exponentially cut-off energy spectrum. the peak energy f£. of which moves out of the X-ray band during the fast decay according to ££.ox£477 with àc11.5.," Their phenomenological model involves an exponentially cut-off energy spectrum, the peak energy $E_{\rm c}$ of which moves out of the X-ray band during the fast decay according to $E_{\rm
+ c}\propto t_{\rm obs}^{-\alpha_{\rm c}}$ with $\alpha_{\rm c}\simeq
+1-1.5$."
+ This fits quite well the present picture. considering in particular that VinXLalobs [Eq. (4))].," This fits quite well the present picture, considering in particular that $\nu_{\rm min}\propto t_{\rm obs}^{-1.5}$ [Eq. \ref{eq:numin}) )]."
+ We also note that some short GRBs show an exponentially decaying light curve around 100s. well beyond the prompt emission. accompanied by spectral evolution. such as GRBOSO724 (Campana et al.," We also note that some short GRBs show an exponentially decaying light curve around 100s, well beyond the prompt emission, accompanied by spectral evolution, such as GRB050724 (Campana et al."
+ 2006). while some show fast decay without apparent late time recovery (e.g. GRBOS1210. GRBOO60801. see Nakar 2007).," 2006), while some show fast decay without apparent late time recovery (e.g. GRB051210, GRB060801, see Nakar 2007)."
+ Regarding the shallow decay phase. a contribution from the reverse shock has been envisaged in Uhm Beloborodov (2007) and Genet et al. (," Regarding the shallow decay phase, a contribution from the reverse shock has been envisaged in Uhm Beloborodov (2007) and Genet et al. ("
+2007). although some tuning appear required to ensure a smooth transition to the recovery phase.,"2007), although some tuning appear required to ensure a smooth transition to the recovery phase."
+ After the present work was submitted. a paper by Petropoulou et al. (," After the present work was submitted, a paper by Petropoulou et al. ("
+2011) appeared. arguing that the shallow decay phase can be accounted for by the low energy tail of the synchrotron self Compton component.,"2011) appeared, arguing that the shallow decay phase can be accounted for by the low energy tail of the synchrotron self Compton component."
+ Alternatively. one could try to explain the shallow decay phase with an inefficient contribution of the forward shock. the efficiency increasing with time and reaching its maximum at recovery of the standard light curve.," Alternatively, one could try to explain the shallow decay phase with an inefficient contribution of the forward shock, the efficiency increasing with time and reaching its maximum at recovery of the standard light curve."
+ This could be accomplished if a small fraction χε With \ lis stored in an accelerated electron powerlaw at that time — beyond the thermal component that amounts to oflxN)oc. — with x rising up to I at recovery., This could be accomplished if a small fraction $\chi\epsilon_e$ with $\chi\ll1$ is stored in an accelerated electron powerlaw at that time – beyond the thermal component that amounts to $\epsilon_e(1-\chi)\sim\epsilon_e$ – with $\chi$ rising up to $\sim1$ at recovery.
+ Granot et al. (, Granot et al. (
+2006) and Ioka et al. (,2006) and Ioka et al. (
+2006) have proposed a similar scenario. with varying microphysical parameters during the shallow decay phase.,"2006) have proposed a similar scenario, with varying microphysical parameters during the shallow decay phase."
+" Our model assumes a sharp transition at Yin.=1 between no acceleration (Le. V= 0) and fully efficient acceleration (4.= 1). but what actually happens at Ying~O(1) when a few efolds of growth of the turbulence canoccur is not known. since we have at our disposal the results of only two simulations at m,=LO+ and 10.7."," Our model assumes a sharp transition at $Y_{\rm inst}=1$ between no acceleration (i.e. $\chi=0$ ) and fully efficient acceleration $\chi=1$ ), but what actually happens at $Y_{\rm inst}\sim {\cal O}(1)$ when a few efolds of growth of the turbulence canoccur is not known, since we have at our disposal the results of only two simulations at $\sigma_{\rm
+ u}=10^{-4}$ and $10^{-5}$."
+ Moreover. one should recall that current PIC simulations probe tiny timescales in regards of the GRB timescales and that these simulations do not vet converge to a stationary shock state (Keshet et al.," Moreover, one should recall that current PIC simulations probe tiny timescales in regards of the GRB timescales and that these simulations do not yet converge to a stationary shock state (Keshet et al."
+ 2009)., 2009).
+ One thus cannot exclude that inefficient acceleration occurs at μοι7O(1) but goes undetected in current simulations: dedicated PIC simulations on long timescales appear required to probe this transition region., One thus cannot exclude that inefficient acceleration occurs at $Y_{\rm inst}\sim {\cal O}(1)$ but goes undetected in current simulations; dedicated PIC simulations on long timescales appear required to probe this transition region.
+ Alternatively. if the jet is structured in energy and Lorentz factor per solid angle. an observer may receive emission from regions of different Lorentz factors than that on the line of sight (e.g. Panaitescu 2007): if the Lorentz factors in those off-axis regions are such that τοιος one might detect low flux emission. corresponding to V.«I and possibly a shallow decay phase.," Alternatively, if the jet is structured in energy and Lorentz factor per solid angle, an observer may receive emission from regions of different Lorentz factors than that on the line of sight (e.g. Panaitescu 2007); if the Lorentz factors in those off-axis regions are such that $Y_{\rm inst}<Y_{\rm c}$, one might detect low flux emission, corresponding to $\chi<1$ and possibly a shallow decay phase."
+ Yet another possibility is that of a clumpy circumburst medium. with clumps of various sizes. provided ro<<rfi5.," Yet another possibility is that of a clumpy circumburst medium, with clumps of various sizes, provided $r_{\rm c}\,\ll\, r/\gamma_{\rm
+ b}$."
+" As the causal region of lateral extent ήτα, contains many clumps. one does not expect a bumpy signature in the light curve."," As the causal region of lateral extent $r/\gamma_{\rm b}$ contains many clumps, one does not expect a bumpy signature in the light curve."
+ However. one would collect only a fraction y.«1 of the expected X-ray flux due to acceleration in the fraction f.«1 of the clumps that carry a magnetization such that Yin«Y. at a given time.," However, one would collect only a fraction $\chi<1$ of the expected X-ray flux due to acceleration in the fraction $f_{\rm c}<1$ of the clumps that carry a magnetization such that $Y_{\rm inst}<Y_{\rm c}$ at a given time."
+ As the overall density and Lorentz factor decrease. f. increases and so does X until recovery. which corresponds to Yi.do in the smallest scale clumps that carry most of the mass.," As the overall density and Lorentz factor decrease, $f_{\rm c}$ increases and so does $\chi$ until recovery, which corresponds to $Y_{\rm
+ inst}\lesssim Y_{\rm c}$ in the smallest scale clumps that carry most of the mass."
+ In each of the above scenarios. one would expect a smooth transition in the light curve with no spectral evolution between the shallow decay phase and the late time normal decay phase. as reported by Nousek et al. (," In each of the above scenarios, one would expect a smooth transition in the light curve with no spectral evolution between the shallow decay phase and the late time normal decay phase, as reported by Nousek et al. ("
+2006) and O'Brien et al. (,2006) and O'Brien et al. (
+2006).,2006).
+ More work is certainly warranted to discuss these aspects in more detail and to compare the properties of the light curve to observational data in the relevant wavelength domains., More work is certainly warranted to discuss these aspects in more detail and to compare the properties of the light curve to observational data in the relevant wavelength domains.
+ One may in particular expect the inverse Compton GeV emission to provide further constraints on the present scenario., One may in particular expect the inverse Compton GeV emission to provide further constraints on the present scenario.
+ we acknowledge support from the CNRS PEPS/PTI Program and from the GDR PCHE., we acknowledge support from the CNRS PEPS/PTI Program and from the GDR PCHE.
+uext periastron passage will take place in July 2011. leaving us plenty of time lor a deep theoretical study of the accretion process before the uext event.,"next periastron passage will take place in July 2014, leaving us plenty of time for a deep theoretical study of the accretion process before the next event."
+ We thank Amit Ixashi for helpful comments., We thank Amit Kashi for helpful comments.
+ This research was supported by the Asher Func for Space Research at the Techinion. aud the Israel Science foundation.," This research was supported by the Asher Fund for Space Research at the Technion, and the Israel Science foundation."
+We refined this transformation locally. in the region of the nova. in the following wav.,"We refined this transformation locally, in the region of the nova, in the following way."
+ We registered the MOSAIC image to the WF3 image using the above transformation., We registered the MOSAIC image to the WF3 image using the above transformation.
+ We then extracted a small subimage surrounding the position of the nova (75x87 WE3 pixels) in ihe MOSAIC and the WF3 image., We then extracted a small subimage surrounding the position of the nova (75x87 WF3 pixels) in the MOSAIC and the WF3 image.
+ The WF3 subimage was convolved with a Gaussian to match the point spread function (PSF) of the MOSAIC image., The WF3 subimage was convolved with a Gaussian to match the point spread function (PSF) of the MOSAIC image.
+ These (wo images were then cross correlated to detect any offset., These two images were then cross correlated to detect any offset.
+ An offset in x of 0.35 px. and in v of -0.70 px. was detected with an accuracy of 0.05 WE2 pixels. which improved the transformation locally by at least a factor of ten.," An offset in x of 0.35 px, and in y of -0.70 px, was detected with an accuracy of 0.05 WF3 pixels, which improved the transformation locally by at least a factor of ten."
+ We then produced a transformation from each of 14 Tenagra images (with the nova well observed) to the MOSAIC image using over 400 stars in common., We then produced a transformation from each of 14 Tenagra images (with the nova well observed) to the MOSAIC image using over 400 stars in common.
+ This produced transformations with a tvpieal RAIS of less than 0.001 MOSAIC pixels., This produced transformations with a typical RMS of less than 0.001 MOSAIC pixels.
+ We tsecl the same cross-correlation technique described above but were unable to detect any local offset in these transformations., We used the same cross-correlation technique described above but were unable to detect any local offset in these transformations.
+ We transformed the nova positions to the MOSAIC coordinate system and then used the MOSAIC to ΑΕΤΟ transformation. plus the detected offsets. to put the nova positions in the WE3 coordinate svstem to better than 0.1 WES pixels.," We transformed the nova positions to the MOSAIC coordinate system and then used the MOSAIC to WF3 transformation, plus the detected offsets, to put the nova positions in the WF3 coordinate system to better than 0.1 WF3 pixels."
+ We (hen used an error weighted average of the 14 positions (o gel a final position in the ΝΕΟ images., We then used an error weighted average of the 14 positions to get a final position in the WF3 images.
+ This position had an RAIS of 0.04 WE3 pixels., This position had an RMS of 0.04 WF3 pixels.
+ Figurigure 7. shows| the region in the WF3 WF3V image surrounding the 1nova. withith a fieldfiek of view of 3 arcseconds on a side.," Figure \ref{m32n1hst} shows the region in the WF3 V image surrounding the nova, with a field of view of 3 arcseconds on a side."
+ We do not convincingly detect the progenitor of M32 nova 1., We do not convincingly detect the progenitor of M32 nova 1.
+ The star indicated by the letter A is at V of 25.5 and I of 24.5. but is well outside the position error circle defined by the RAIS of the nova positions.," The star indicated by the letter A is at V of 25.5 and I of 24.5, but is well outside the position error circle defined by the RMS of the nova positions."
+ We estimate that the progenitor of this nova was fainter than V. of 26.0 at the epoch that these data were taken., We estimate that the progenitor of this nova was fainter than V of 26.0 at the epoch that these data were taken.
+ This gives an amplitude of at least 8.7 magnitudes in V for (he nova outburst., This gives an amplitude of at least 8.7 magnitudes in V for the nova outburst.
+ The WEPC2 ] and F1042M images were well registered with the V image and nothing was detected in them at the position of the nova., The WFPC2 I and F1042M images were well registered with the V image and nothing was detected in them at the position of the nova.
+ We conclude. therefore. that M32 nova 1 could not be a," We conclude, therefore, that M32 nova 1 could not be a"
+In Figure 2 we show the comparison of the stellar-vs-halo mass relation for the BCGs in our AGN-ON/OFF simulations with the predictions of abundance matching (?) (black solid line).,In Figure \ref{fig:AM} we show the comparison of the stellar-vs-halo mass relation for the BCGs in our AGN-ON/OFF simulations with the predictions of abundance matching \citep{Moster:2010p5423} (black solid line).
+ The BCG mass predicted by the AGN-OFF model deviates more than 3c from the average relation obtained through abundance matching., The BCG mass predicted by the AGN-OFF model deviates more than $3\sigma$ from the average relation obtained through abundance matching.
+" Without feedback, the BCG forms stars at an efficiency of from the cosmic available baryon fraction, thus it lies close to the dotted red line."," Without feedback, the BCG forms stars at an efficiency of from the cosmic available baryon fraction, thus it lies close to the dotted red line."
+" On the contrary, the AGN-ON model prediction is remarkably close to the average relation, and well inside the lo scatter bars."," On the contrary, the AGN-ON model prediction is remarkably close to the average relation, and well inside the $1\sigma$ scatter bars."
+" For comparison with recent simulations of massive early-type galaxies that do not include AGN feedback, we show the results of ? and ?.."," For comparison with recent simulations of massive early-type galaxies that do not include AGN feedback, we show the results of \cite{Naab:2009p5731} and \cite{Feldmann:2010p1516}."
+ The mass of the central group galaxy simulated by these authors are both ~30 above the average relation; also in this case the result is close to the simplified model with star formation efficiency., The mass of the central group galaxy simulated by these authors are both $\sim3\sigma$ above the average relation; also in this case the result is close to the simplified model with star formation efficiency.
+ This suggests that AGN feedback could well be the mechanism that reduces the star formation efficiencies in massive galaxies., This suggests that AGN feedback could well be the mechanism that reduces the star formation efficiencies in massive galaxies.
+ We constructed mock images of the central region of the simulated galaxy clusters to detect further differences between the two models., We constructed mock images of the central region of the simulated galaxy clusters to detect further differences between the two models.
+" Figure 4 shows a set of such images at z=1 and z=0, where the color of each pixel represents the flux in the g’ band."," Figure \ref{fig:gbandmaps} shows a set of such images at $z=1$ and $z=0$, where the color of each pixel represents the flux in the $g'$ band."
+ We have chosen the same spatial and color scales in all the pictures to show the remarkable difference in sizes between the AGN-ON and OFF simulations., We have chosen the same spatial and color scales in all the pictures to show the remarkable difference in sizes between the AGN-ON and OFF simulations.
+" The AGN-OFF BCG is much more extended than the AGN-ON one, at z=1 as well as at z=0, ie. the global size of the galaxies scales with their mass."," The AGN-OFF BCG is much more extended than the AGN-ON one, at $z=1$ as well as at $z=0$, i.e. the global size of the galaxies scales with their mass."
+" However, the halflight isophotes (black contours) show that the stellar light is more concentrated in the AGN-OFF BCG than in the AGN-ON one."," However, the half-light isophotes (black contours) show that the stellar light is more concentrated in the AGN-OFF BCG than in the AGN-ON one."
+ This implies that the surface density profile of the AGN-ON galaxy is shallow when compared to the AGN-OFF one; this is indeed the case as we will show in 3.4.., This implies that the surface density profile of the AGN-ON galaxy is shallow when compared to the AGN-OFF one; this is indeed the case as we will show in \ref{subsec:profiles}.
+" To estimate the range of possible ellipticities an observer would see, we view the simulations from different anglesand make the same surface brightness images."," To estimate the range of possible ellipticities an observer would see, we view the simulations from different anglesand make the same surface brightness images."
+" We then fit the half-light isophotes with ellipses and estimate their ellipticity €=1—b/a, where b and a are the semi-minor and semi-major axis of the ellipses."," We then fit the half-light isophotes with ellipses and estimate their ellipticity $\epsilon=1-b/a$, where $b$ and $a$ are the semi-minor and semi-major axis of the ellipses."
+ In this way we give the range of possible ellipticities for our galaxies., In this way we give the range of possible ellipticities for our galaxies.
+" At z=0 the AGN-ON BCG has e=[0.123, 0.155], while the AGN-OFF BCG has e—[0.257,0.498]; both ranges are consistent with values typical of spheroidal systems, but the AGN-ON galaxy is slightly more spherical."," At $z=0$ the AGN-ON BCG has $\epsilon=[0.123,0.155]$ , while the AGN-OFF BCG has $\epsilon=[0.257,0.498]$; both ranges are consistent with values typical of spheroidal systems, but the AGN-ON galaxy is slightly more spherical."
+" At z—1 we have e=[0.201,0.424] for the AGN-ON BCG, and €=[0.120,0.721] for the AGN-OFF BCG; at this redshift, the AGN-ON galaxy is slightly more spheroidal than at z= 0, while the ellipticity of the AGN-OFFgalaxy is more typical of flattened ellipsoidal systems."," At $z=1$ we have $\epsilon=[0.201,0.424]$ for the AGN-ON BCG, and $\epsilon=[0.120,0.721]$ for the AGN-OFF BCG; at this redshift, the AGN-ON galaxy is slightly more spheroidal than at $z=0$ , while the ellipticity of the AGN-OFFgalaxy is more typical of flattened ellipsoidal systems."
+ Figure 5 shows the average line of sight velocity maps, Figure \ref{fig:velmaps} shows the average line of sight velocity maps
+"The non-conservative potential of the bar could also produce an increment of the energy of the halo particles and these could reach the escape velocity producing some kind of ""evaporation"" of the halo.",The non-conservative potential of the bar could also produce an increment of the energy of the halo particles and these could reach the escape velocity producing some kind of “evaporation” of the halo.
+ The mechanism ts the same as the one explained in previous subsections but the initial conditions will be different from those explained in §2.1.5.., The mechanism is the same as the one explained in previous subsections but the initial conditions will be different from those explained in \ref{.initial}.
+ Indeed. the fact that we get some net accretion of some IGM particles arises because some of them go through the centre of the galaxy. but the virtalized halo has totally different initial conditions.," Indeed, the fact that we get some net accretion of some IGM particles arises because some of them go through the centre of the galaxy, but the virialized halo has totally different initial conditions."
+ Assuming that the halo particles have an initial velocity corresponding to a circular velocity plus isotropic dispersion velocities with a component in the radial direction [given the isotropy. the dispersion in other directions will be the same: the isotropic representation. while more or less correct in the inner halo. is not correct in the outer halo (Dehnen McLaughlin 2005). but we consider it here as a rough approximation] Y=35/18 (Dehnen MeLaughlin 2005). in agreement with simulations and analytical solutions of self-gravitating dynamical equilibriums; with the halo density of eq. (11).," Assuming that the halo particles have an initial velocity corresponding to a circular velocity plus isotropic dispersion velocities with a component in the radial direction [given the isotropy, the dispersion in other directions will be the same; the isotropic representation, while more or less correct in the inner halo, is not correct in the outer halo (Dehnen McLaughlin 2005), but we consider it here as a rough approximation] $\alpha =35/18$ (Dehnen McLaughlin 2005), in agreement with simulations and analytical solutions of self-gravitating dynamical equilibriums; with the halo density of eq. \ref{halodens}) ),"
+ and 77=a)vGO/V3 (Dehnen MeLaughlin 2005). where νο) is the circular velocity. we can perform again the numerical calculation of the orbits and the increment of energy according to eqs. (4))," and $\sigma _r (r=a)=v_c(a)/\sqrt{3}$ (Dehnen McLaughlin 2005), where $v_c(r)$ is the circular velocity, we can perform again the numerical calculation of the orbits and the increment of energy according to eqs. \ref{deltaE}) )"
+ and (600)., and \ref{Fphi}) ).
+ The result with 100000 halo particles whose initial conditions are distributed randomly according to the probability densities of positions and velocities including the dispersions. during 14. Gyr orbit calculation. gives the result that none of these particles gets an increase in energy AE>ee where ro is the initial distance from the centre of the galaxy.," The result with 000 halo particles whose initial conditions are distributed randomly according to the probability densities of positions and velocities including the dispersions, during 14 Gyr orbit calculation, gives the result that none of these particles gets an increase in energy $\Delta E>\frac{GM(r_0)}{2r_0}$, where $r_0$ is the initial distance from the centre of the galaxy."
+ This means that the increase in the halo evaporation due to the bar-halo interaction is negligible: +M<~107+ during the whole life of the galaxy., This means that the increase in the halo evaporation due to the bar–halo interaction is negligible: $\frac{\Delta M}{M}<\sim 10^{-4}$ during the whole life of the galaxy.
+ Let us suppose that the dependence on the initial angle. initial velocity and the halo mass are separable.," Let us suppose that the dependence on the initial angle, initial velocity and the halo mass are separable."
+ This is not strictly correct. but may be appropriate as a rough estimation.," This is not strictly correct, but may be appropriate as a rough estimation."
+ In this case. a generalization of the expression (31)) will be: Here we are assuming that the dise and the bar do not change with time.," In this case, a generalization of the expression \ref{form200}) ) will be: Here we are assuming that the disc and the bar do not change with time."
+" The solution of this differential equation will be: and the total accretion in the lifetime of the Galaxy (we take r=1410"" yr) due to this mechanism is The necessary conditions for a fraction F of the halo mass to be accreted by this mechanism (0. AMΞFMpato(t): assuming there are no other mechanisms of accretion| are: Particularly. for Mya(r)=12κ107 Μ.. &,=45° and vo50 kni/s. This expression has sense only for p(kg/m?)<4.7x1077 Καπ (F< 1). for higher densities the mass of the halo should be larger."," The solution of this differential equation will be: and the total accretion in the lifetime of the Galaxy (we take $t=1.4\times 10^{10}$ yr) due to this mechanism is The necessary conditions for a fraction $F$ of the halo mass to be accreted by this mechanism $\Delta M=FM_{\rm halo}(t)$ ; assuming there are no other mechanisms of accretion] are: Particularly, for $M_{\rm halo}(t)=1.2\times 10^{12}$ $_\odot $, $\theta _\infty =45^\circ $ and $v_\infty=50$ km/s, This expression has sense only for $\rho _\infty ({\rm kg/m^3})\le 4.7\times
+10^{-23}$ $^3$ $F\le 1$ ), for higher densities the mass of the halo should be larger."
+ In Fig. 8..," In Fig. \ref{Fig:Fdens},"
+ the values of F for different densities are plotted., the values of $F$ for different densities are plotted.
+ Densities of the [IGM around 1077+ ke/m* (including baryonie and non-baryonic matter) are able to produce a significant amount of accretion by means of this mechanism., Densities of the IGM around $10^{-24}$ $^3$ (including baryonic and non-baryonic matter) are able to produce a significant amount of accretion by means of this mechanism.
+" In order to have at least of the total halo matter acereted by this mechanism (F= 0.1) it is necessary that (we keep &,,=45°. ΛΜ=1.2«107 Mj) The expression (41)) for F>0.1 indicates that for ος245°. Μο)=1.2κ107 Ma. v=50 km/s we need a density of 2x1077 kg/m?23x10 Μ./ΜΡΟΣ=10°57ρω."," In order to have at least of the total halo matter accreted by this mechanism $F=0.1$ ) it is necessary that (we keep $\theta _\infty =45^\circ $, $M_{\rm halo}(t)=1.2\times 10^{12}$ $_\odot $ ) The expression \ref{F1}) ) for $F>0.1$ indicates that for $\theta _\infty =45^\circ $, $M_{\rm halo}(t)=1.2\times 10^{12}$ $_\odot $, $v_\infty =50$ km/s we need a density of $2\times 10^{-24}$ $^3=3\times 10^{13}$ $_\odot $ $^3=10^2h_0^{-2}\rho _{\rm critical}$."
+ This is a high density and may be reached by the intracluster gas in some groups or clusters of galaxies. especially in the richest ones. which can have a density much higher than this (including dark matter).," This is a high density and may be reached by the intracluster gas in some groups or clusters of galaxies, especially in the richest ones, which can have a density much higher than this (including dark matter)."
+ The typical IGM densities in clusters are between 10? and 10-7 kg/m? (Roussel et al., The typical IGM densities in clusters are between $10^{-25}$ and $10^{-23}$ $^3$ (Roussel et al.
+ 2000)., 2000).
+ However.the relativevelocities in such environments are also higher. thus decreasing the desired effect (see the next subsection).," However,the relativevelocities in such environments are also higher, thus decreasing the desired effect (see the next subsection)."
+ For cases such as the Local Group this density 1s too high., For cases such as the Local Group this density is too high.
+ The, The
+Ap=0.091 mag (Schlegeletal.1998).. adopting (he extinction curve in Cardelli.andMathis (1989).,"$A_B = 0.091$ mag \citep{sch98}, adopting the extinction curve in \citet{car89}."
+.. The 9 spectroscopically confirmed clusters wilh A-band data are shown with filled circles., The 9 spectroscopically confirmed clusters with $K$ -band data are shown with filled circles.
+ Also shown are single-aged stellar population (SSP) models by for ages of 2. 3. 4. 5. 6. 8. LL and 14 Gyrs and metallicities of |Z/II]|2 —2.25. —1.35. —0.33. 0.0 and 470.35.," Also shown are single-aged stellar population (SSP) models by \citet{mar02} for ages of 2, 3, 4, 5, 6, 8, 11 and 14 Gyrs and metallicities of = $-2.25$, $-1.35$, $-0.33$, 0.0 and $+0.35$."
+ From the optical colors alone ancl assuming old ages. the 9 clusters would appear to have low to intermediate metallicities. but. their location in the two-color diagram suggests that thev may instead be as voung as 23 Gyirs and have metallicities near solar.," From the optical colors alone and assuming old ages, the 9 clusters would appear to have low to intermediate metallicities, but their location in the two-color diagram suggests that they may instead be as young as 2–3 Gyrs and have metallicities near solar."
+ For a more detailed discussion of the(VA... —7)) diagram. including acomparison of various SSP models. we refer to D02.," For a more detailed discussion of the, ) diagram, including acomparison of various SSP models, we refer to P02."
+ In Fig., In Fig.
+ 7 we show the Dalmer line indices (I12. II54 and H94) for the clusters aand(Fe). compared with SSP model grids by Thomas.Maraston. TAIBO2)..," \ref{fig:idxs} we show the Balmer line indices $\beta$, $\gamma_A$ and $\delta_A$ ) for the clusters and, compared with SSP model grids by \citet[][hereafter TMB02]{tmb02}."
+ A variety of SSP models are now available in the literature. but [or this work we have chosen the models by TMDO2 because they are tabulated for several different vvalues and (hus allow us to quantify the effect of varving a-element abundances.," A variety of SSP models are now available in the literature, but for this work we have chosen the models by TMB02 because they are tabulated for several different values and thus allow us to quantify the effect of varying $\alpha$ -element abundances."
+ These models are also the first to attempt acorrection for the bbias in the original Lick/IDS fitting functions (Marastonetal.2002)., These models are also the first to attempt acorrection for the bias in the original Lick/IDS fitting functions \citep{marea02}.
+. We have used the common definitions == (Fe5270+Fe5335)/2 and == \/Algh(Fe) (González1993).," We have used the common definitions = $+$ Fe5335)/2 and = $\sqrt{{\rm Mg}b\,\fe}$ \citep{gon93}."
+. The SSP models are shown for ages of 3. 5. 8. 11 and 14 Gvrs and the same metallicities as in Fig. 3..," The SSP models are shown for ages of 3, 5, 8, 11 and 14 Gyrs and the same metallicities as in Fig. \ref{fig:pvk_vi}."
+ In the plots involving HL9. we show models with (o/Fe]=0 (solid lines) and [a/Fe|=+0.5 (dashed lines).," In the plots involving $\beta$, we show models with $\alphafe=0$ (solid lines) and $\alphafe=+0.5$ (dashed lines)."
+ In the other plots only [a/Fe]=0 models ave shown., In the other plots only $\alphafe=0$ models are shown.
+ Chisters wil and without A-band imaging are shown with filled and open circles. respectively.," Clusters with and without $K$ -band imaging are shown with filled and open circles, respectively."
+ Object #33. which has Mgb «0. is shown at [MgFe|—0.," Object 3, which has $b$ $<0$, is shown at $\mgfe=0$."
+ The blue color of this cluster. as well as ils weak Fe indices. suggest that it is verv metal-poor.," The blue color of this cluster, as well as its weak Fe indices, suggest that it is very metal-poor."
+" It is also one of the faintest objects in our sample and was observed in a short (6"") slit. making accurate sky subtraction diffieult."," It is also one of the faintest objects in our sample and was observed in a short $6\arcsec$ ) slit, making accurate sky subtraction difficult."
+ The negative Mgb value is most likely due to contamination by the [INI] night-skv line at 5199A. which is shifted into the Mgb central bandpass after correcting for the radial velocity of NGC 4365.," The negative $b$ value is most likely due to contamination by the ] night-sky line at 5199, which is shifted into the $b$ central bandpass after correcting for the radial velocity of NGC 4365."
+ All low panels in Fig., All four panels in Fig.
+ 7 confirm a spread in metallicity with some clusters approaching Solar values. but none of the clusters included in our small saaiple appear to have metallicities significantly above Solar.," \ref{fig:idxs} confirm a spread in metallicity with some clusters approaching Solar values, but none of the clusters included in our small sample appear to have metallicities significantly above Solar."
+ The combined iidex measures primarily Fe-peak elements (Tripicco&Bell1995) aud is therefore sensitive to wwhen plotted for a fixed mean metallicity [7/1]., The combined index measures primarily Fe-peak elements \citep{tb95} and is therefore sensitive to when plotted for a fixed mean metallicity .
+ The eerids. on the other hand. vary only weakly with [ον fixed |[Z/1I]. as can be seen from the similarity between the dashed-line and solid-line," The grids, on the other hand, vary only weakly with for fixed , as can be seen from the similarity between the dashed-line and solid-line"
+ (Jamultonotal.1991:Weiuberg (Naravanan Drenieretal.2003:Mohavaee2006) (Lavauxotal.2010).. 2011)..," \citep{ham91,wei92,pea94,pea96,fri02,bre03,moh06}, \citep{nar99} \citep{fri02,bre03,moh06} \citep{lav10}. ."
+ The displacement field is a crucial quantity in most of these methods. aud is of interest in its own rieht.," The displacement field is a crucial quantity in most of these methods, and is of interest in its own right."
+ There is reason to believe that a logaritlianic raustormmation of the density field may aid estinatiou of the displacement field., There is reason to believe that a logarithmic transformation of the density field may aid estimation of the displacement field.
+ The mass distribution has (een successfullv described by a lognormal field that evolved from Gaussian initial conditions (Coles&Jones 1991).. with some evidence for a skewed loguorma field iu the nonlinear regime (Colomibi1991).," The mass distribution has been successfully described by a lognormal field that evolved from Gaussian initial conditions \citep{coles91}, with some evidence for a skewed lognormal field in the nonlinear regime \citep{col94}."
+. More recently. it has been found that the power spectrum of he log-transformed density Geld coutains more Fisher information than the usual power spectrum at ual scales by up to a factor of —10 (Nevrincketal.2009).," More recently, it has been found that the power spectrum of the log-transformed density field contains more Fisher information than the usual power spectrum at small scales by up to a factor of $\sim$ 10 \citep{ney09}."
+. A modified logarithmic transform has also been shown Oo increase the precision of the power spectrum of he nonlinear weal lensing convergence feld (Seoctal. 2010)., A modified logarithmic transform has also been shown to increase the precision of the power spectrum of the nonlinear weak lensing convergence field \citep{seo10}.
+. The log-transformed deusity field is also nore effective iu coustraiiue cosmological parameters han the standard density field when using the power spectra (Nevrinek2011)., The log-transformed density field is also more effective in constraining cosmological parameters than the standard density field when using the power spectrum \citep{ney11}.
+. This paper investigates the effect of a logaritlianic ranstorm of the deusity field ou the relation between tle density aud displacement fields., This paper investigates the effect of a logarithmic transform of the density field on the relation between the density and displacement fields.
+ We do this by 1icasuriug he divergence of the displacement field. which iu lucar heory is proportional to the negative density contrast à. using both Eulerian aud Lagrangian techniques iu a coswnological N-body simulation.," We do this by measuring the divergence of the displacement field, which in linear theory is proportional to the negative density contrast $\delta$, using both Eulerian and Lagrangian techniques in a cosmological $N$ -body simulation."
+ We compare the incar and logarithmic approaches aud evaluate their dependence on redshift and smoothing scale. both of which affect the applicability of linear theory.," We compare the linear and logarithmic approaches and evaluate their dependence on redshift and smoothing scale, both of which affect the applicability of linear theory."
+ The linear and logarithmic approxinations forthe displacement field are derived iu Section 2.. aud three methods of," The linear and logarithmic approximations forthe displacement field are derived in Section \ref{sec:reconstruction}, , and three methods of"
+"cosmology,",cosmology.
+ We have then adapted their model to our prelerrecl cosmology. keeping all other parameters the same.," We have then adapted their model to our preferred cosmology, keeping all other parameters the same."
+ We have confirmed (hat in (he new cosmology. the derived. blazar Iuninosity function reproduces (he EGRET 2nd catalog flux distribution of observed blazars. in exactly the same wav as in the original Stecker-Salamon calculation.," We have confirmed that in the new cosmology, the derived blazar luminosity function reproduces the EGRET 2nd catalog flux distribution of observed blazars, in exactly the same way as in the original Stecker-Salamon calculation."
+" In the case of normal egalaxies. /44,4,ln can be expressed in terms of the CSER function (mass being converted to stars per unit time per unit comoving volume. denoted by 5,(z))."," In the case of normal galaxies, $\ngdot$ can be expressed in terms of the CSFR function (mass being converted to stars per unit time per unit comoving volume, denoted by $\dot{\rho}_\star(z)$ )."
+ The star formation rate (SFR. mass converted to stars per unit time) of an individual galaxy ∖∖↽↕∐∣↽≻≼↲≼⇂≼↲∐∪∥↲≼⇂∣↽≻∡∖↽," The star formation rate (SFR, mass converted to stars per unit time) of an individual galaxy will be denoted by $\psi$."
+∣⇁⋅↕∐∪↕⋅≺⇂≼↲↕⋅↥∪≀↧⊍∖⊽⋟∖⇁∪≺⊲↕≀↧↴∥↲⊔∐↲≼↕⊲⊳↔⊲∏⋧∖∖↽↕⊔↥∣∎∣⋅ νοt we assume that: (1) the ∐↕↖≺↽↔↴∐↕↥↓≀↧⊔∖⋱∖⊽≼↲↕∐⇂∪↓⋟↥∐↲↕∐∐↕≀↧↴↥∐↕≀↧⊔∖⊽⋟∖⇁↓⋟∏∐≺∢∐∪∐⊔↼∖↕⊡↕⋟∖⊽∏↕∐∖⇁≼↲↕⋱∖⊽≀↧↴⋅≀↧↴↕∐⇂⊔⋯⋟∖⊽∣⇁≀↧⊍∖⇁≼⇂≼↲≺⇂∏≺∢≼↲≼⊔⋟↕⋅∪∐↕ observations of high-nmass stars is always proportional to the supernova explosion rate in the same galaxy: (2) the cosmic rav [lux in a galaxy is proportional to. ( and (he cosmic ray spectral shape is universal (see Fieldsefaf (2001))): and (2) at any cosmic epoch the cosmic rav escape properties are (he same as in (he present Milkv. Way. and any 5-rays. produced aller escape are negligible.," In order to associate the CSFR with $\ngdot$, we assume that: (1) the high mass end of the initial mass function (IMF) is universal, and thus $\psi$ as deduced from observations of high-mass stars is always proportional to the supernova explosion rate in the same galaxy; (2) the cosmic ray flux in a galaxy is proportional to $\psi$ and the cosmic ray spectral shape is universal (see \citet{focv}) ); and (3) at any cosmic epoch the cosmic ray escape properties are the same as in the present Milky Way, and any $\gamma$ -rays produced after escape are negligible."
+ An average galaxvs ?-rav. (nimber) luminosity is. by virtue of our assumptions. ο ο ∖∖⊽∐≼↲↕⋅≼↲⊡↕⊳∖⊽⊔∐↲↕↽≻∐∪↥∪∐≼↲∐≼↲↕⋅≸≟⋡∖↽↕∐⊔∐↲≸↽↔↴≀↧↴↥��↧↴⇀↸∡∖⋰⊳∖⇁↕⋅≼↲⊳∖⊽↥∐⋅≀↧↴∐∐↲⋅≀↧↴∐≼⇂∕∣⋜⋰∶⇄⋝↕⊳∖⊽⊔∐↲↖⊂↽↔↴≀↧↪∖⇁∐↓≀↧↪∖⊽⊳∖⇁ ↓⋟↕⋅≀↧↴≺∢∐∪∐∪↓⋟⊔∐↲≸↽↔↴≀↧↴↥≀∟↸⋡∖↽≀↧↴↥↕⋅≼↲≼⇂⊳∖⇁∐↕∐∠⋅↴∏∐↲↓≯≀↧↴≺∢↥∪↕⋅∕∣≼⋝⊳∶↕⋝∕∕∕∕∣≼⋝∩⇄⋝↥⋯↪∖⊽∣↽≻≼↲≼↲∐↥∐∏⋅∪⋔∐∢≼↲≼⇂↥∪⋯∢≺∢∪∏∐↥ ↓≯∪↕⋅⊔∐↲↕∐≺∢↕⋅≼↲≀↧↪∖⊽≼↲∪↓⋟↥≀↧↴↕⋅≸≟≼↲↥≀↧↴↥∪∐↓⊳∖⊽≀," An average galaxy's $\gamma$ -ray (number) luminosity is, by virtue of our assumptions, (z,E) = (E) where $E$ is the photon energy in the galaxy's rest frame, and $\mu(z)$ is the gas mass fraction of the galaxy at redshift z. The factor $\mu(z)/\mu(0)$ has been introduced to account for the increase of target atoms at earlier cosmic epochs and assumes a “closed box” galaxy."
+"↧↴↥≼↲≀↧↴↕⋅∐≼↲↕⋅≺∢∪⊳∖⊽∐↓↕≺∢≼↲↕↽≻∪≺∢∐∖⇁≀↧↴↕∐⇂≀↧↪∖⊽⊳∖⇁∏∐∐↲⋝∖⊽≀↧↴⋅⋅≺∢↥∪⊳∖⇁≼↲≼⊓↽≻∪⇀↸⋮⋅≸↽↔↴≀↧↴↥≀↧↴⇀↸⋡∖↽⋅ The emissivity density will then be eon((2.£)⋀Ym --m ο.D(4) where the comoving galaxy number density. 7,4 and star formation rate c(z) are related to the CSER via tna=f."," The emissivity density will then be (z,E)= = (E), where the comoving galaxy number density $n_{\rm gal}$ and star formation rate $\psi(z)$ are related to the CSFR via $\psi n_{\rm gal} = \dot{\rho}_\star$."
+ Note that. due to our assumptions. the conversion of a certain amount of gas into stars will result to the production of the same amount of 5 ravs from CIR-ISM interactionsgalaxies.," Note that, due to our assumptions, the conversion of a certain amount of gas into stars will result to the production of the same amount of $\gamma$ rays from CR-ISM interactions."
+ Therefore. our calculation does not depend on the observational knowledge of both i and Nea Individually. but only on that of their product. which is the CSER.," Therefore, our calculation does not depend on the observational knowledge of both $\psi$ and $n_{\rm gal}$ individually, but only on that of their product, which is the CSFR."
+ Now the sum of gas mass and star mass in a closed box is constant in (ime. equal to the barvonic mass of the galaxy.," Now the sum of gas mass and star mass in a closed box is constant in time, equal to the baryonic mass of the galaxy."
+ Hence.," Hence,"
+1979).,.
+. Coronal spectra. used for the incident N-rayv. racliation fiekl were computed. using emissivities from the CIILANTI database and the ion populations of (1998).. as implemented in the IDL suite of programs 2000).," Coronal spectra used for the incident X-ray radiation field were computed using emissivities from the CHIANTI database and the ion populations of , as implemented in the IDL suite of programs ."
+. For a given coronal N-rav specirum. the maximum intensitv of a [luorescent line is achieved lor a heliocentric angle 8=0 and height /=0.," For a given coronal X-ray spectrum, the maximum intensity of a fluorescent line is achieved for a heliocentric angle $\theta=0$ and height $h=0$."
+ As in the case of the earlier study ol Ne. since we are primarily interested here in (he observability of the fluorescent line. we adopt 9=0 and /=0 as baseline parameters. these values vielding the maxinunm possible line strength.," As in the case of the earlier study of Ne, since we are primarily interested here in the observability of the fluorescent line, we adopt $\theta=0$ and $h=0$ as baseline parameters, these values yielding the maximum possible line strength."
+ The wavelength of the O Ika doublet has not been determined with high precision iil we adopted the value 23.37-(£0.02A.. corresponding to the mean of values found for the energy. level of the excited [15|2p!!P of 530.6+0.3 eV [rom laboratory measurements bv(1994).. and2005)mainDocdvyCitationEnd934|Stolte.," The wavelength of the O $\alpha$ doublet has not been determined with high precision and we adopted the value $23.37\pm 0.02$, corresponding to the mean of values found for the energy level of the excited $[1s]2p^4\,^4P$ of $530.6\pm 0.3$ eV from laboratory measurements by, and."
+etal:97. This wavelength is in good agreement with the position of the 2p'!S—[1s|2p!!P. resonance seen in absorption in the interstellar medium toward bright N-rav continuum sources by(2004).," This wavelength is in good agreement with the position of the $2p^3\,^4S \rightarrow
+[1s]2p^4\,^4P$ resonance seen in absorption in the interstellar medium toward bright X-ray continuum sources by."
+". The O ko flux was computed for isothermal irradiating coronal spectra. will plasma temperatures in the range 10""-10* IX and coronal and photospheric chemical compositions of andO/II28.", The O $\alpha$ flux was computed for isothermal irradiating coronal spectra with plasma temperatures in the range $10^6$ $10^7$ K and coronal and photospheric chemical compositions of and.
+66).. As noted in 811. these two chemical mixtures differ not onlv in O abundance but also. primarily. in C. N and Ne abundances: (hese differences (urn out to be important for the O fluorescent line and are discussed in more detail below.," As noted in 1, these two chemical mixtures differ not only in O abundance but also, primarily, in C, N and Ne abundances: these differences turn out to be important for the O fluorescent line and are discussed in more detail below."
+ To invesügate changes in (he O abundance only. we also performed caleulations using a photospherie composition corresponding to the GS abundance mixture with O elevated and decreased by a factor of 2 (O/11—9.13.8.53. or (O/1I]—20.3).," To investigate changes in the O abundance only, we also performed calculations using a photospheric composition corresponding to the GS abundance mixture with O elevated and decreased by a factor of 2 (O/H=9.13,8.53, or $\pm0.3$ )."
+ The photospheric O Ίνα fIuorescent line is shown in comparison with the direct coronal spectrum [or the case of GS abundances for both coronal spectrum. aud photosphere in, The photospheric O $\alpha$ fluorescent line is shown in comparison with the direct coronal spectrum for the case of GS abundances for both coronal spectrum and photosphere in
+pseudorapidity distribution of charged particles in inelastic events.,pseudorapidity distribution of charged particles in inelastic events.
+" In order to calculate the dN.,/dn for inelastic events we summed charged particles pseudorapidity distributions of single-diffractive and non-single diffractive events with the default theoretically calculated weights (solid lines) [7] as well as with the weights deduced from the experimental data (dotted lines) [2, 14].."," In order to calculate the $dN_{ch}/d\eta$ for inelastic events we summed charged particles pseudorapidity distributions of single-diffractive and non-single diffractive events with the default theoretically calculated weights (solid lines) \cite{HMD} as well as with the weights deduced from the experimental data (dotted lines) \cite{UA5diff, UA5PRP}. ."
+" We calculate dN/dn using the following expression: As we see, the predictions of QGSM for inelastic evens are significantly different from experimental data in 7 region from 1 to 3, whereas the predictions for single-diffractive and non-single diffractive events are described well At V/s = 900 GeV pp single-diffractive interactions UAS5 triggers were not able to register particles produced from diffractedsystems with masses below 2.5 GeV/c?. "," We calculate $dN/d\eta$ using the following expression: As we see, the predictions of QGSM for inelastic evens are significantly different from experimental data in $\eta$ region from 1 to 3, whereas the predictions for single-diffractive and non-single diffractive events are described well At $\sqrt{s}$ = 900 GeV $p\bar{p}$ single-diffractive interactions UA5 triggers were not able to register particles produced from diffractedsystems with masses below 2.5 $c^2$ "
+a single magnetic element observed at increasingly better resolutions should induce a dramatic increase in the polarimetric signal: while the latter because à microturbulent scenario is not compatible with the extended tails that evidence a cascade of spatial scales coexisting in the internetwork quiet Sun.,a single magnetic element observed at increasingly better resolutions should induce a dramatic increase in the polarimetric signal; while the latter because a microturbulent scenario is not compatible with the extended tails that evidence a cascade of spatial scales coexisting in the internetwork quiet Sun.
+ The correct interpretation of the quiet Sun should take into account that the magnetic field vectorB depends on the scale r and the description has to rely on the complete AsensioRamos 2009b)..," The correct interpretation of the quiet Sun should take into account that the magnetic field vector$\mathbf{B}$ depends on the scale $r$ and the description has to rely on the complete scale-dependent continuous probability distribution function $p(\mathbf{B}_1,r_1;\mathbf{B}_2,r_2;\ldots;\mathbf{B}_n,r_n)$ \citep[e.g.,][]{andres09}."
+ If the scale of oreanization is much smaller than the resolution element. a large degree of cancellations occurs and the distribution of observed polarization amplitudes quickly tends to a Gaussian distribution.," If the scale of organization is much smaller than the resolution element, a large degree of cancellations occurs and the distribution of observed polarization amplitudes quickly tends to a Gaussian distribution."
+" Therefore, the Gaussian core of the circular polarization amplitude histogram can be identified with the fields organized well below the resolution element."," Therefore, the Gaussian core of the circular polarization amplitude histogram can be identified with the fields organized well below the resolution element."
+ The extended tails are thus identified with magnetic elements organized at a larger scale., The extended tails are thus identified with magnetic elements organized at a larger scale.
+" We calculate a “characteristic size"" of the smallest magnetic elements through inversion with a model atmosphere that is formed, in each resolution element, by an isotropic distribution of magnetic field vectors all of them of the same size."," We calculate a “characteristic size” of the smallest magnetic elements through inversion with a model atmosphere that is formed, in each resolution element, by an isotropic distribution of magnetic field vectors all of them of the same size."
+" With this model we get an upper limit to the size of these magnetic elements —10 km, which is smaller than the photon mean free path in the photosphere."," With this model we get an upper limit to the size of these magnetic elements $\sim$ 10 km, which is smaller than the photon mean free path in the photosphere."
+ Such small scale structuring of the quiet Sun magnetism was conjectured by Sánchez(1996) based on the ubiquity of the Stokes asymmetries., Such small scale structuring of the quiet Sun magnetism was conjectured by \cite{jorge_egidio_valentin_96} based on the ubiquity of the Stokes asymmetries.
+"However, although the majority of magnetic fields in the quiet Sun is organized at such small scales, models taking into account the cascade of scales are necessary to understand the quiet Sun magnetism.","However, although the majority of magnetic fields in the quiet Sun is organized at such small scales, models taking into account the cascade of scales are necessary to understand the quiet Sun magnetism."
+" The direct analysis of observables reveals that there is not a substantial change in the polarization amplitudes from 1.3"" to 0.32"".", The direct analysis of observables reveals that there is not a substantial change in the polarization amplitudes from $''$ to $0.32''$.
+ A fraction of of the circular (linear) signals follow a Gaussian (Rayleigh) distribution and seem to be insensitive to the spatial resolution., A fraction of of the circular (linear) signals follow a Gaussian (Rayleigh) distribution and seem to be insensitive to the spatial resolution.
+ This important fact reveals that the Zeeman effect is indeed sensitive to the microturbulent magnetic field usually diagnosed with the Hanle effect., This important fact reveals that the Zeeman effect is indeed sensitive to the microturbulent magnetic field usually diagnosed with the Hanle effect.
+" In fact. even a thousand of mixed-polarity, magnetic elements in the resolution element gives a detectable Zeeman signal 2007).."," In fact, even a thousand of mixed-polarity, magnetic elements in the resolution element gives a detectable Zeeman signal \citep{arturo07}. ."
+ The remainder of the observed polarization signals show a modification with, The remainder of the observed polarization signals show a modification with
+truncations do not necessarily occur at the same ancl galactocentric distances or with the same abruptness on either side (e.g. INSI. Jensen Thuan 1982. Nàsslund Jérrsatter 1997. Abe et al.,"Disc truncations do not necessarily occur at the same galactocentric distances or with the same abruptness on either side (e.g., KS1, Jensen Thuan 1982, Nässlund Jörrsätter 1997, Abe et al."
+ 1999. Fry et al.," 1999, Fry et al."
+ 1999)., 1999).
+ In. most cases. however. the truncation radii on either side of the ealactic disc occur within ~10154 of cach other.," In most cases, however, the truncation radii on either side of the galactic disc occur within $\sim 10-15$ of each other."
+ The profiles in Figs., The profiles in Figs.
+" 5. and 6.. and our estimates for Buys in ‘Tablelable 3 stshow that in general.&L the eldisces of fouour samplele galgalaxies are truncated at similar ασ, within their observational uncertainties. with the possible exception of ESO 201-Ci22."," \ref{cutoffs.fig} and \ref{Iband.fig}, and our estimates for $R_{\rm max}$ in Table \ref{Rmax.tab} show that in general, the discs of our sample galaxies are truncated at similar radii, within their observational uncertainties, with the possible exception of ESO 201-G22."
+ Table 3 also contains our best estimates for the exponentia scalelength. in the truncation region. Pis.," Table \ref{Rmax.tab} also contains our best estimates for the exponential scalelength in the truncation region, $h_{R,\delta}$."
+ To measure this scalelength. we defined the inner fitting radius as the radius where the radial surface brightness profile starts to deviate significantly from the model radial exponential light. distribution (οἱ.," To measure this scalelength, we defined the inner fitting radius as the radius where the radial surface brightness profile starts to deviate significantly from the model radial exponential light distribution (cf."
+ Barteldrees Dettmar 1994): the uncertainty in this radius is generally 1554.. depending on the particular galaxy considered.," Barteldrees Dettmar 1994); the uncertainty in this radius is generally $\le 15$, depending on the particular galaxy considered."
+ As outer fitting boundary we used Zeus., As outer fitting boundary we used $R_{\rm max}$.
+ Phe errors associated. with fas are observational uncertainties. obtained from the comparison of several similar fits in which we adjusted the inner boundary of the radial fitting range by20'4.," The errors associated with $h_{R,\delta}$ are observational uncertainties, obtained from the comparison of several similar fits in which we adjusted the inner boundary of the radial fitting range by."
+. The uncertainties in {τμ are included in the observational errors given: the formal errors were l.. The ellects on disc truncations and. scale parameters of an inclination jzonQ are negligible for our galaxy sample of ;>ST galaxies. as was convincingly shown by Barteldrees & (1994). in particular in view of the systematic Disc observational uncertainties involved.," The uncertainties in $R_{\rm max}$ are included in the observational errors given; the formal errors were $\le 1$ The effects on disc truncations and scale parameters of an inclination $i \neq 90^\circ$ are negligible for our galaxy sample of $i \ge 87^\circ$ inclined galaxies, as was convincingly shown by Barteldrees Dettmar (1994), in particular in view of the systematic and observational uncertainties involved."
+ The combination of fps and 9 gives us an indication of the asymmetry and sharpness of the actual disc truncations.," The combination of $h_{R,\delta}$ and $\delta$ gives us an indication of the asymmetry and sharpness of the actual disc truncations."
+ Considering. the relatively large systematic and observational errors that cannot be avoided. at this point. we cannot claim that we detect any systeniatic asvnimetries. perhaps with the exception of ESO 416-Ci25.," Considering the relatively large systematic and observational errors that cannot be avoided at this point, we cannot claim that we detect any systematic asymmetries, perhaps with the exception of ESO 416-G25."
+ The northern edge of the disc of ESO 416-€Ci25 is very sharply runcated compared to its southern edge., The northern edge of the disc of ESO 416-G25 is very sharply truncated compared to its southern edge.
+ Upon close examination of the actual CCD images. we believe that this may be explained by the fact that we likely observe the outer stellar envelope of a spiral arm. whereas on the southern side we are looking into the inside of a spiral arm.," Upon close examination of the actual CCD images, we believe that this may be explained by the fact that we likely observe the outer stellar envelope of a spiral arm, whereas on the southern side we are looking into the inside of a spiral arm."
+ Note. however. that also at the southern edge of the disc a clear runcaion signature is observed (Eig. 5))," Note, however, that also at the southern edge of the disc a clear truncation signature is observed (Fig. \ref{cutoffs.fig}) )."
+ The last two columns of Table 3.2 show that in none of our sample galaxies the radial scalelength in the truncation region decreases to values of order or less than 1 kpe., The last two columns of Table \ref{Rmax.tab} show that in none of our sample galaxies the radial scalelength in the truncation region decreases to values of order or less than 1 kpc.
+ Using Hm=50 km s 1. or other values closer. to the current best estimate. will further increase the truncation scalcloneths measured in our galaxies.," Using $H_0 = 50$ km $^{-1}$ $^{-1}$, or other values closer to the current best estimate, will further increase the truncation scalelengths measured in our galaxies."
+" Although two of our sample galaxies may not be exactly edge-on. their inclinations are sulliciently close to 907 so as not to increase the measurements of 255,5 bx more than their observational uncertainties (ef"," Although two of our sample galaxies may not be exactly edge-on, their inclinations are sufficiently close to $^\circ$ so as not to increase the measurements of $h_{R,\delta}$ by more than their observational uncertainties (cf."
+Darteldrees Dettmar 1904).,Barteldrees Dettmar 1994).
+ We are therefore forced to conclude that. although our clises are," We are therefore forced to conclude that, although our discs are"
+a lobe-dominated: double source. both the spectral index and the position of the optical source strongly support this hypothesis.,"a lobe-dominated double source, both the spectral index and the position of the optical source strongly support this hypothesis."
+ There are various facts that support the conclusion that RX J1011.215545 is an AGN. the most relevant being a hieh racio to optical flux ratio. à 210keV luminosity exceeding 10eres+ and the broad. Mell emission.," There are various facts that support the conclusion that RX J1011.2+5545 is an AGN, the most relevant being a high radio to optical flux ratio, a $2-10\, \keV$ luminosity exceeding $10^{45}\, \ulum$ and the broad MgII emission."
+ Ln addition the strong NeV|A3624 line implies the presence of an underlying iard ionising continuum., In addition the strong $\lambda$ 3624 line implies the presence of an underlying hard ionising continuum.
+ We initially suspected obscuration in this source rccause ofits high. PSPC hardness ratio., We initially suspected obscuration in this source because of its high PSPC hardness ratio.
+ Also for a typical uncovered AGN. he average optical magnitude corresponding to its N-ray lux would be A—19 (see. e.g. Hasinger 1996) instead of he observed value of /?—21.," Also for a typical uncovered AGN, the average optical magnitude corresponding to its X-ray flux would be $R\sim 19$ (see, e.g., Hasinger 1996) instead of the observed value of $R\sim 21$."
+ Lhe absence of a broad. CLL] ine confirms this obscuration hypothesis., The absence of a broad CIII] line confirms this obscuration hypothesis.
+ A weak broad Mgll line is detected. its equivalent width xing 3 to 5 times smaller than for a tvpe lL AGN (Francis et al 1991: Baker Llunstead 1995).," A weak broad MgII line is detected, its equivalent width being 3 to 5 times smaller than for a type I AGN (Francis et al 1991; Baker Hunstead 1995)."
+ his cannot be explained as a simple obscuration effect. since in that case both the broad lines and he nuclear continuum would be equally suppressed. leaving he equivalent widths unchanged.," This cannot be explained as a simple obscuration effect, since in that case both the broad lines and the nuclear continuum would be equally suppressed, leaving the equivalent widths unchanged."
+ Ehe weakness of Mell and he absence of broad. CIV. and Hell may be the result. o£ dilution by a source of blue continuum over and above that emanating cirectIy from the nucleus., The weakness of MgII and the absence of broad CIV and HeII may be the result of dilution by a source of blue continuum over and above that emanating directly from the nucleus.
+ Ehe requirement would oe that at à rest wavelength of ~2800 the nuclear continuum may be only 20 to 50 per cent of the otal., The requirement would be that at a rest wavelength of $\sim 2800$ the nuclear continuum may be only 20 to 50 per cent of the total.
+ The nature of this extra blue component is unknown. out. rellected nuclear radiation. nebular continuum and copious star formation are all possibilities.," The nature of this extra blue component is unknown, but reflected nuclear radiation, nebular continuum and copious star formation are all possibilities."
+ The detection of a rellected Fe Ix. line in N-ravs and the strong OLY line with respect to typical type LE situation favour the enhanced star formation scenario., The non-detection of a reflected Fe K line in X-rays and the strong [OII] line with respect to typical type I situation favour the enhanced star formation scenario.
+ The equivalent width of the broad. CLL component is expected to be roughly 2 to 5 times smaller than that of Megll in a tvpe L AGN. and therefore it would be very weak in this object.," The equivalent width of the broad CIII] component is expected to be roughly 2 to 5 times smaller than that of MgII in a type I AGN, and therefore it would be very weak in this object."
+ Obscuration of the nuclear continuum could also lead to the narrow Nev] lines having enhanced equivalent widths., Obscuration of the nuclear continuum could also lead to the narrow [NeV] lines having enhanced equivalent widths.
+ The power-law in the X-ray spectrum of this object is similar to that. found. for other luminous raclio-loud quasars at high. redshifts. (ον1.5. Cappi et al 1997). distinctively Hatter than for radio-quiet ACGN.," The power-law in the X-ray spectrum of this object is similar to that found for other luminous radio-loud quasars at high redshifts, $\Gamma\sim 1.5$, Cappi et al 1997), distinctively flatter than for radio-quiet AGN."
+ This has been associated with dillerent emission mechanisms (svnchrotron sclf-Compton with the radio-emitting electrons in racio-Loud AGN versus nuclear emission in racio-quict objects), This has been associated with different emission mechanisms (synchrotron self-Compton with the radio-emitting electrons in radio-loud AGN versus nuclear emission in radio-quiet objects).
+ lt ids then possible that in racio-loud active galaxies the line-ol-sight to N-rav. emitting regions intercepts less obscuring material than does the direct. path to the nucleus., It is then possible that in radio-loud active galaxies the line-of-sight to X-ray emitting regions intercepts less obscuring material than does the direct path to the nucleus.
+ Larger absorbing columns (Ng~1077em7) than that observed in RN J1011.215545 are common only among radio-loud quasars at very high redshifts (23. Cappi et al 1997. Fiore et al 1998).," Larger absorbing columns $N_H\sim
+10^{22}\, \ucol$ ) than that observed in RX J1011.2+5545 are common only among radio-loud quasars at very high redshifts $z>3$, Cappi et al 1997, Fiore et al 1998)."
+ Phe possible contribution to the X-ray Lux from a cluster of galaxies hosting this source (which might be dominant in radiogalaxies. Crawford Fabian 1996) is small. since the X-ray. data does not show evidence for a spectral eutolf consistent with thermal emission.," The possible contribution to the X-ray flux from a cluster of galaxies hosting this source (which might be dominant in radiogalaxies, Crawford Fabian 1996) is small, since the X-ray data does not show evidence for a spectral cutoff consistent with thermal emission."
+ The amount of X-ray absorption predicts an optical extinction for the X-ray source which is ly=1.1ji. using stancard dust to gas ratios.," The amount of X-ray absorption predicts an optical extinction for the X-ray source which is $A_V=1.1^{+6.7}_{-0.85}$, using standard dust to gas ratios."
+ For moderate extinction 1 2). the nuclear light seen in the optical can be direct radiation from the nucleus.," For moderate extinction $A_V\sim 1-2$ ), the nuclear light seen in the optical can be direct radiation from the nucleus."
+ However. if the obscuration is much Larger. then the Mgll broad line would be seen through rellection only.," However, if the obscuration is much larger, then the MgII broad line would be seen through reflection only."
+ Lt is even. possible that the nucleus is very heavily obscured in the optical (il 10) in which case the direct N-rav. continuum and nuclear Fe Ix. emission. might also be suppressed. leaving a dominant. N-ray component arising in the radio lobes with only moderate associated photoclectric absorption.," It is even possible that the nucleus is very heavily obscured in the optical $A_V\gg 10$ ) in which case the direct X-ray continuum and nuclear Fe K emission might also be suppressed, leaving a dominant X-ray component arising in the radio lobes with only moderate associated photoelectric absorption."
+ Disentaneling both possibilities requires high spatial resolution optical and LR observations., Disentangling both possibilities requires high spatial resolution optical and IR observations.
+ In any event. the discovery of this object demonstrates that high-redshift) racio-loucd obscured AGN are. present ab faint X-ray dluxes.," In any event, the discovery of this object demonstrates that high-redshift radio-loud obscured AGN are present at faint X-ray fluxes."
+ Such objects may ῥίαν a role. albeit probably minor. in producing the X-ray. background.," Such objects may play a role, albeit probably minor, in producing the X-ray background."
+ Surveys to be carried out with ΑΝΛΙ and. XNMM will undoubtely find large numbers of obscured AGNs ancl show what is their contribution to the X-ray backerouncl., Surveys to be carried out with AXAF and XMM will undoubtely find large numbers of obscured AGNs and show what is their contribution to the X-ray background.
+" XD and RC were visiting astronomers of the Centro-Astronómnmico LHispano-Memánn. Calar Alto. operated. by the Max-Plancek-bIastitute for Astronomy. Heidelberg jointly with the Spanish ""Comisiónn Nacional de ia."," XB and RC were visiting astronomers of the Centro-Astron\'ommico Hispano-Alemánn, Calar Alto, operated by the Max-Planck-Institute for Astronomy, Heidelberg jointly with the Spanish `Comisiónn Nacional de a'."
+ Phe William Llersehel Telescope is operated. on. the island. of La Palma by the Isaac Newton Group in the spanish Observatorio cel Roque de los Muchachos of the Instituto de, The William Herschel Telescope is operated on the island of La Palma by the Isaac Newton Group in the spanish Observatorio del Roque de los Muchachos of the Instituto de
+Figure 7 compares aco as a function of metallicity between this study and the literature.,Figure \ref{fig:fancy_plot} compares $\alpha_{\rm CO}$ as a function of metallicity between this study and the literature.
+ The points in red indicate IR-based measurements., The points in red indicate IR-based measurements.
+" In detail, our measurements (circles) yieldlower aco than previous IR-based studies."," In detail, our measurements (circles) yield $\alpha_{\rm CO}$ than previous IR-based studies."
+ We suspect that this is mainly because we solve for aco without assuming óapg or measuring it far away from the region of interest., We suspect that this is mainly because we solve for $\alpha_{\rm CO}$ without assuming $\delta_{\rm GDR}$ or measuring it far away from the region of interest.
+" One likely sense of systematic variations in Óapg is that dapr is likely to be higher in the dense gas close to molecular complexes, which tend to reside mainly in the stellar disk, than in a diffuse, extended ddisk (e.g.,???).."," One likely sense of systematic variations in $\delta_{\rm GDR}$ is that $\delta_{\rm GDR}$ is likely to be higher in the dense gas close to molecular complexes, which tend to reside mainly in the stellar disk, than in a diffuse, extended disk \citep[e.g.,][]{STANIMIROVIC99,DRAINE07B,MUNOZMATEOS09}."
+" If óapm is taken to be too high, Equation 2 yields a corresponding overestimate of aco."," If $\delta_{\rm GDR}$ is taken to be too high, Equation \ref{eq:model}
+ yields a corresponding overestimate of $\alpha_{\rm CO}$."
+" In the SMC our attempt to remove a diffuse component along the line of sight also leads to lower aco 5.1)), though it is less clear that our approach is (Sectioncorrect in that case."," In the SMC our attempt to remove a diffuse component along the line of sight also leads to lower $\alpha_{\rm CO}$ (Section \ref{sec:comments}) ), though it is less clear that our approach is correct in that case."
+" Regardless of the cause, by simultaneously solving foraco and óapn in the regions"," Regardless of the cause, by simultaneously solving for$\alpha_{\rm CO}$ and $\delta_{\rm GDR}$ in the regions"
+I1IC4N derived by (Wxyrowskietal.2002).,$_3$ N derived by \citep{wyr2002}.
+. To fit the IICN lines we use the temperature. density and velocity structure of the expanding envelope. which fits the IIC4N lines. and only vary the IICN abundance.," To fit the HCN lines we use the temperature, density and velocity structure of the expanding envelope, which fits the $_3$ N lines, and only vary the HCN abundance."
+ Using a temperature of 560 Ix. (see discussion below) a ICN column density of 2x10PIS 7 is needed to cause the observed absorption., Using a temperature of 560 K (see discussion below) a HCN column density of $2\times 10^{18}$ $^{-2}$ is needed to cause the observed absorption.
+ The resulting fils to the spectra are shown in Figure 7 together with a spectrum of the IIC4N ο=1 J=12-11 transition which has a similar upper energv (L300 KIN) Since no continuum {lux measurements were performed al 350 Gllz. the HON J=4—3 spectra are shown in brightness temperature units. whereas for (he other spectra the ratio of line to continuum temperature. which reduces calibration uncertainties. is shown.," The resulting fits to the spectra are shown in Figure \ref{3d_model} together with a spectrum of the $_3$ N $v_4=1$ $J=12-11$ transition which has a similar upper energy $\sim 1300$ K) Since no continuum flux measurements were performed at 350 GHz, the HCN $J=4-3$ spectra are shown in brightness temperature units, whereas for the other spectra the ratio of line to continuum temperature, which reduces calibration uncertainties, is shown."
+ The deviation of the ICN J—43 model from the observed spectrum could be due to pointing and/or focus errors: a pointing error of aalone would explain the difference between observation and model aud cannot be excluded., The deviation of the HCN $J=4-3$ model from the observed spectrum could be due to pointing and/or focus errors: a pointing error of alone would explain the difference between observation and model and cannot be excluded.
+ The model consists of power laws for temperature. density and velocity starüng al an inner radius of amd an region within that radius. which fits the continuum measurements of CRL 618.," The model consists of power laws for temperature, density and velocity starting at an inner radius of and an region within that radius, which fits the continuum measurements of CRL 618."
+ The temperature al the inner radius is 560 Ix. In the model. the IICN J=4—3 line is highly oplically Chick aud mostly sensitive to the model temperature and (he emitting size.," The temperature at the inner radius is 560 K. In the model, the HCN $J=4-3$ line is highly optically thick and mostly sensitive to the model temperature and the emitting size."
+ The ]1ICN direct. (type lines. on the other hand. are optically thin ancl probe a combinationof temperature and column density. of (he model.," The HCN direct $\ell$ -type lines, on the other hand, are optically thin and probe a combinationof temperature and column density of the model."
+ The best fit model has a HCN/IICSN abundance ratio of 3 to 6 dependent on the assumed population of the vibrational levels of 1IC4N which is consistent with the result obtained by mid-infrared absorption measurements by Cernicharoetal.(2001)., The best fit model has a $_3$ N abundance ratio of 3 to 6 dependent on the assumed population of the vibrational levels of $_3$ N which is consistent with the result obtained by mid-infrared absorption measurements by \citet{cer2001}.
+. Figure 8. shows the results of the VLA observations., Figure \ref{vla-hcn} shows the results of the VLA observations.
+ The total continuum flux density al 40 GIIz. estimated from the (lux on the shortest baselines. is 0.75 Jv with an uncertainty ol1054.," The total continuum flux density at 40 GHz, estimated from the flux on the shortest baselines, is 0.75 Jy with an uncertainty of."
+".. The size of the continuum eniission is 0.34x0.16"".. estimated [rom a Gaussian fit to the UV data."," The size of the continuum emission is $0.34\times 0.16$, estimated from a Gaussian fit to the UV data."
+" To increase the spectral sensitivity. every four channels were averaged together and a taper in the UV plane was applied. reducing the angular resolution to 0.34x0.31""."," To increase the spectral sensitivity, every four channels were averaged together and a taper in the UV plane was applied, reducing the angular resolution to $0.34\times 0.31$."
+ The insert in Figure 5. shows the spectrum integrated over the indicated area., The insert in Figure \ref{vla-hcn} shows the spectrum integrated over the indicated area.
+ Line parameters of a Gaussian fit to the spectrum are given in Table 2.., Line parameters of a Gaussian fit to the spectrum are given in Table \ref{analysis}.
+ To image the HCN absorption the continuum was subtracted from the data using the channel ranges marked in the insert of Figure 8.., To image the HCN absorption the continuum was subtracted from the data using the channel ranges marked in the insert of Figure \ref{vla-hcn}.
+ The contours in Figure 8. show (he LCN absorption averaged over the line., The contours in Figure \ref{vla-hcn} show the HCN absorption averaged over the line.
+ The observed 40 GlIlz contünuun emission compares well with the results of (1993.1995). at 23 GIIz.," The observed 40 GHz continuum emission compares well with the results of \citet{martin-pintado93,martin-pintado95} at 23 GHz."
+ The HCN absorption falls into the same velocity range as (he hot core (IC) component seen in ammonia by (1992).., The HCN absorption falls into the same velocity range as the hot core (HC) component seen in ammonia by \citet{mar92}. .
+ The IICN absorption is slightly shifted to the west from the center of the continuum.," The HCN absorption is slightly shifted to the west from the center of the continuum,"
+in this study are comparable in magnitude to those presented by Getman (20082.b) and Aarnio (2010) but significantly smaller than (hose presented by Favata (2005).,"in this study are comparable in magnitude to those presented by Getman (2008a,b) and Aarnio (2010) but significantly smaller than those presented by Favata (2005)."
+ This discrepancy is most probably caused by the much longer observation (ime of sources in the COUP mission than in our sample. which allows for (the discovery of more extended flares with longer decay times.," This discrepancy is most probably caused by the much longer observation time of sources in the COUP mission than in our sample, which allows for the discovery of more extended flares with longer decay times."
+" Comparison of the [lares in our sample with a well-studied superllare on II Pegasi. a inagnelically active main sequence star. reveals (hat although the peak flare temperature Class Land LE YSOs. and the lack of loop lengths longer than 13x10!"" em on the identified Class HII YSOs. suggests and energy emitted bv the II Peg superllare is comparable to energies in our sample. (he entire II Peg superflare event was relatively short."," Comparison of the flares in our sample with a well-studied superflare on II Pegasi, a magnetically active main sequence star, reveals that although the peak flare temperature Class I and II YSOs, and the lack of loop lengths longer than $13
+\times 10^{10}$ cm on the identified Class III YSOs, suggests and energy emitted by the II Peg superflare is comparable to energies in our sample, the entire II Peg superflare event was relatively short."
+ Overall. the the II Pee superllare is consistent with the compact flares in our sample. suggesting analagous [lare production mechanisms.," Overall, the the II Peg superflare is consistent with the compact flares in our sample, suggesting analagous flare production mechanisms."
+ A search for 6.4 keV fhiorescent iron emission as observed in the I] Peg superílare in our sample vielded a null result., A search for 6.4 keV fluorescent iron emission as observed in the II Peg superflare in our sample yielded a null result.
+ This was however. consistent with the number of stars in our sample and (he generally observed low rate of detection.," This was however, consistent with the number of stars in our sample and the generally observed low rate of detection."
+ The detection of flares of order several stellar radii in length on the existence of long magnetic structures connecting the star and disk., The detection of flares of order several stellar radii in length on the existence of long magnetic structures connecting the star and disk.
+ We suggest the following scenario for magnetic interaction: equatoriallv located magnetic loops like those seen on (hie Sun can erow large enough to interact with the inner edge of a cireumstellar accretion disk. provided convection on the chromosphere does not shuttle the loop footprint (oo much.," We suggest the following scenario for star-disk magnetic interaction: equatorially located magnetic loops like those seen on the Sun can grow large enough to interact with the inner edge of a circumstellar accretion disk, provided convection on the chromosphere does not shuffle the loop footprint too much."
+ The presence of a circumstellar disk does not determine flare energeties. but should a field line reach out and [ind disk material. the disk supports this extended magnetic structure.," The presence of a circumstellar disk does not determine flare energetics, but should a field line reach out and find disk material, the disk supports this extended magnetic structure."
+ When the magnetic loop ruptures due to. e... differential rotation between the star and the disk. an extended flare results.," When the magnetic loop ruptures due to, e.g., differential rotation between the star and the disk, an extended flare results."
+ Star-clisk magnetic coupling thus explains the distribution of loop lengtlis seen in Figure 4: the thin or non-existent disks associated with Class Η YSOs cannot support the coherent magnetic field structures required for long Mares., Star-disk magnetic coupling thus explains the distribution of loop lengths seen in Figure 4: the thin or non-existent disks associated with Class III YSOs cannot support the coherent magnetic field structures required for long flares.
+ A limitation of the hycrodyvnamic model emploved in this study is the inability to distinguish between a single large flare loop or an arcade of smaller loops lor a given loop length L. However. we find it eviclentiary (hat the only flares with L of order the dust destruction radius occur on stars wilh disks. aud we do not find extended flare loop lengths for Class HII sources.," A limitation of the hydrodynamic model employed in this study is the inability to distinguish between a single large flare loop or an arcade of smaller loops for a given loop length L. However, we find it evidentiary that the only flares with L of order the dust destruction radius occur on stars with disks, and we do not find extended flare loop lengths for Class III sources."
+ While these results are nol conclusive proof of star-disk magnetic interaction. they are hiehlv suggestive.," While these results are not conclusive proof of star-disk magnetic interaction, they are highly suggestive."
+ The authors gratefully acknowledge many useful comments from the anonvimous referee. as well as the assistance of Ettore Flaccomio with X-ray data reduction. and the assistance ol Robert Gutermuth in sharing IRAC and MIPS data for the sources in our sample.," The authors gratefully acknowledge many useful comments from the anonymous referee, as well as the assistance of Ettore Flaccomio with X-ray data reduction, and the assistance of Robert Gutermuth in sharing IRAC and MIPS data for the sources in our sample."
+ This publication makes use of data products from the Two Micron All Sky Survey. which is a joint project of the University of Massachusetts and the Infrared Processing and Analvsis Center. [funded by the National Aeronauties anc Space Administration and the National Science," This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center, funded by the National Aeronautics and Space Administration and the National Science"
+calibrated spectra of the source.,calibrated spectra of the source.
+ We then analyzed them with the GILDAS-CLASS90software!., We then analyzed them with the GILDAS-CLASS90.
+. The single side band spectrum of receiver Ib was obtained through a standard manual procedure., The single side band spectrum of receiver 1b was obtained through a standard manual procedure.
+ We first removed spurious features (one) and then compared all scans one by one for a given frequency to remove unwanted lines from the corresponding image side band by blanking out the corresponding channels., We first removed spurious features (one) and then compared all scans one by one for a given frequency to remove unwanted lines from the corresponding image side band by blanking out the corresponding channels.
+" In the few cases where image band lines were blended with signal band lines, a fit was performed to the blended lines to separate the contribution from each band."," In the few cases where image band lines were blended with signal band lines, a fit was performed to the blended lines to separate the contribution from each band."
+ The image band emission was then subtracted from the signal band spectrum., The image band emission was then subtracted from the signal band spectrum.
+ This procedure was repeated twice in order to ensure a complete cleaning of lines from the image side band at each frequency setting., This procedure was repeated twice in order to ensure a complete cleaning of lines from the image side band at each frequency setting.
+" Once the deconvolution was done for one of the receivers, the other receiver was treated automatically by blanking out the same frequency blocks in each individual scan."," Once the deconvolution was done for one of the receivers, the other receiver was treated automatically by blanking out the same frequency blocks in each individual scan."
+ Linear baselines were subtracted and all 1456 single side band individual scans (728 per receiver) merged., Linear baselines were subtracted and all 1456 single side band individual scans (728 per receiver) merged.
+" The calibration and the double side band gain ratio was checked against the strongest lines, and found to be consistent within2-4%."," The calibration and the double side band gain ratio was checked against the strongest lines, and found to be consistent within."
+. The final spectrum with a spectral resolution of kmss! is shown in 11., The final spectrum with a spectral resolution of $^{-1}$ is shown in 1.
+" While ground-based observations have provided much insight into the chemical inventory, the unprecedented spectral coverage of HIFI will yield many lines arising from a wide variety of species that together will greatly expand our understanding of the density, temperature, and dynamical structure of the ejecta and the processes driving the molecular complexity of these chemical ""smokestacks"" of the galaxy, from the dust forming zones to the radical and photodominated outer regions."," While ground-based observations have provided much insight into the chemical inventory, the unprecedented spectral coverage of HIFI will yield many lines arising from a wide variety of species that together will greatly expand our understanding of the density, temperature, and dynamical structure of the ejecta and the processes driving the molecular complexity of these chemical “smokestacks” of the galaxy, from the dust forming zones to the radical and photodominated outer regions."
+ In this work we present the first high spectral resolution line survey towards an evolved star in the submillimeter and far-IR domains., In this work we present the first high spectral resolution line survey towards an evolved star in the submillimeter and far–IR domains.
+" Figure11 shows around 130 well-detected lines with intensities exceeding mmK. All of them, except one, can be easily assigned to lines of CO (including v=1; see Patel et al."," 1 shows around 130 well-detected lines with intensities exceeding mK. All of them, except one, can be easily assigned to lines of CO (including $\nu$ =1; see Patel et al."
+" 2009b), HCN and HCN (in several vibrational states), SiO, SiS, and CS (including its isotopologs and vibrationally excited levels), AIF,"," 2009b), HCN and $^{13}$ CN (in several vibrational states), SiO, SiS, and CS (including its isotopologs and vibrationally excited levels), AlF,"
+complicated by the fact that SAIC X.1: was also present in the INTE PCA and HENTE feld-of-view in many of our observations.,complicated by the fact that SMC X–1 was also present in the RXTE PCA and HEXTE field-of-view in many of our observations.
+ Using power spectra [rom each observation. we determined in which observations 31-5 pulsations from NTIE JO111.27317 were present and O.7-s pulsations [rom SALC Xo] were not present.," Using power spectra from each observation, we determined in which observations 31-s pulsations from XTE J0111.2–7317 were present and 0.7-s pulsations from SMC X–1 were not present."
+ We then generated energy spectra anc response matrices for the PCX and LIENTE data for these 20 observations and fitted them in NSPEC with an absorbe power-law with a high energy cut-olf and a Ciaussian iron line., We then generated energy spectra and response matrices for the PCA and HEXTE data for these 20 observations and fitted them in XSPEC with an absorbed power-law with a high energy cut-off and a Gaussian iron line.
+ Averageὃν parameter values for the 20 observations are eiven in Table 1.., Average parameter values for the 20 observations are given in Table \ref{tab:xr_spectra}.
+ Phe flux in the iron line was correlate with the total lux. with a correlation coellicient of 0.07 anc a chance probability of 2.5107. suggesting that the iron line is intrinsic to NTE JO111.21317.," The flux in the iron line was correlated with the total flux, with a correlation coefficient of 0.97 and a chance probability of $2.5 \times 10^{-5}$, suggesting that the iron line is intrinsic to XTE J0111.2–7317."
+ No other parameters were obviously correlated with ας nor did they show clear evolution with time., No other parameters were obviously correlated with flux nor did they show clear evolution with time.
+ Using only those observations where SAIC X-1.: was not present in the power spectrum. we estimated. pulse fractions from RATE PCA data.," Using only those observations where SMC X-1 was not present in the power spectrum, we estimated pulse fractions from RXTE PCA data."
+ Figure 6 shows the pulse fraction versus the 2-50 keV flux computed from our NSPEC fits., Figure \ref{fig:xr_pfrac} shows the pulse fraction versus the 2-50 keV flux computed from our XSPEC fits.
+ The pulse fraction is correlated. with the total Hux. indicating that the increase in (ux is primarily in the pulsed component.," The pulse fraction is correlated with the total flux, indicating that the increase in flux is primarily in the pulsed component."
+ Comparing BATSE and PCA pulsed Huxes. we see evidence wt the first peak of the outburst is harder than the sccone »eak. shown in Figure 7..," Comparing BATSE and PCA pulsed fluxes, we see evidence that the first peak of the outburst is harder than the second peak, shown in Figure \ref{fig:xr_hr}."
+ Unfortunately. because SAIC X- is present in all of the observations during the first. peak of 10 outburst. we cannot study the ellects of this change in iardness on the shape of the energy spectrum in detail.," Unfortunately, because SMC X-1 is present in all of the observations during the first peak of the outburst, we cannot study the effects of this change in hardness on the shape of the energy spectrum in detail."
+ The ohvsical explanation of this is unclear., The physical explanation of this is unclear.
+ We speculate tha js may be an obscuration or absorption elfect - the softer X-rays are partially absorbed or blocked by the large amoun of material present at the beginning of the outburst., We speculate that this may be an obscuration or absorption effect - the softer X-rays are partially absorbed or blocked by the large amount of material present at the beginning of the outburst.
+ As the outburst progresses. this material begins to dissipate an 16 softer N-rays become more easily seen.," As the outburst progresses, this material begins to dissipate and the softer X-rays become more easily seen."
+ A comparison of HENTE and PCA pulsed Iuxes suggests a similar change in hardness. although the LENTE fluxes are much less certain than the BATSE fluxes due to the smaller effective area anc integration time.," A comparison of HEXTE and PCA pulsed fluxes suggests a similar change in hardness, although the HEXTE fluxes are much less certain than the BATSE fluxes due to the smaller effective area and integration time."
+ This suggests that the elfect is intrinsic to 10 source and is most likely not instrumental., This suggests that the effect is intrinsic to the source and is most likely not instrumental.
+ Observations of steady spin-up during the outburs allow us to compute a lower limit to the luminosity of NTTIS J0111.2.7317., Observations of steady spin-up during the outburst allow us to compute a lower limit to the luminosity of XTE J0111.2–7317.
+ Phe angular momentum of a rotating neutron star is given by (=2arlv where v is the spin [requency and / is the moment of inertia of the neutron star., The angular momentum of a rotating neutron star is given by $\ell = 2 \pi I \nu$ where $\nu$ is the spin frequency and $I$ is the moment of inertia of the neutron star.
+ The torque on the neutron star is obtained by. dilferentiating (or N=f2zxIp., The torque on the neutron star is obtained by differentiating $\ell$ or $N = \dot \ell = 2 \pi I \dot \nu$.
+ In an acercline pulsar undergoing steady spin-up. the characteristic torque is given by Nina= where AM is the mass accretion rate. € is the eravitational constant. A/x is the mass of the neutron star.," In an accreting pulsar undergoing steady spin-up, the characteristic torque is given by $N_{\rm max} = \dot M (G M_{\rm X} r_{\rm m})^{1/2}$ where $\dot M$ is the mass accretion rate, $G$ is the gravitational constant, $M_{\rm X}$ is the mass of the neutron star."
+ and αμ is the magnetospheric radius., and $r_{\rm m}$ is the magnetospheric radius.
+ The maximum possible torque occurs when the magnetospheric radius equals the corotation radius., The maximum possible torque occurs when the magnetospheric radius equals the corotation radius.
+ Setting NoNias vields an expression for Al and hence Lx that depends only on the spin frequency and its derivative for assumed. values of the neutron star parameters. 1.6. Assuming /=107 © cni. Mx=LAM... and Ry = 10 kni. Lx>107 [ον typical observed. values. of 5-324 nmllzand £25.10 ls ft.," Setting $N\leq N_{\rm max}$ yields an expression for $\dot M$ and hence $L_{\rm X}$ that depends only on the spin frequency and its derivative for assumed values of the neutron star parameters, i.e. Assuming $I = 10^{45}$ g $^{2}$, $M_{\rm X} = 1.4 M_{\odot}$ , and $R_{\rm X}$ = 10 km, $L_{\rm X} \geq 10^{38}$ for typical observed values of $\nu
+= 32.4$ mHz and $\dot \nu = 5 \times 10^{-11}$ Hz $^{-1}$."
+" The average 20-50 keV. pulsed. ux was 2.510I"" implying wt NTE 011191317 is at a distance of dz59.8 kpe. confirming that it is in the SAIC."," The average 20-50 keV pulsed flux was $2.5 \times 10^{-10}$, implying that XTE J0111.2–7317 is at a distance of $d \geq 59.8$ kpc, confirming that it is in the SMC."
+ The outhurst from NTE J0111.27317. shoulel be Jassified as a giant or Type HL outburst given the Large spin- rates (3.1)«10 Hogs+ measured with BATSE and ie laree luminosities (2—9)107 nmieasured in the PCA observations where SAIC A-1 was not. present.," The outburst from XTE J0111.2–7317 should be classified as a giant or Type II outburst given the large spin-up rates $(3-7) \times 10^{-11}$ Hz $^{-1}$ measured with BATSE and the large luminosities $(2-9)
+\times 10^{38}$, measured in the PCA observations where SMC X-1 was not present."
+ Llowever. the temporal profile of this outhurst is. unlike eint outbursts in other BesX-ray binaries that typically. consist of either a single large intensity peak or a large peak followed by smaller peaks modulated at the orbital period.," However, the temporal profile of this outburst is unlike giant outbursts in other Be/X-ray binaries that typically consist of either a single large intensity peak or a large peak followed by smaller peaks modulated at the orbital period."
+ The initial large peak is usually much brighter and of much longer duration that any succeeding peaks., The initial large peak is usually much brighter and of much longer duration that any succeeding peaks.
+ NTL 011127317 insteacl has two peaks that are similar in intensity with the first peak lasting about 25 days and the second lasting about GO days. followed by a declining tail lasting about 30 days.," XTE J0111.2--7317 instead has two peaks that are similar in intensity with the first peak lasting about 25 days and the second lasting about 60 days, followed by a declining tail lasting about 30 days."
+ The separation between the two peaks is about 30 davs., The separation between the two peaks is about 30 days.
+ “Phis double peaked structure is most likely not due to orbital effects because the frequency derivative is strongly correlated with the pulsed Dux ancl also shows the double peaked structure., This double peaked structure is most likely not due to orbital effects because the frequency derivative is strongly correlated with the pulsed flux and also shows the double peaked structure.
+ Such a correlation is predicted by simple accretion theory for svstems accreting from a disk., Such a correlation is predicted by simple accretion theory for systems accreting from a disk.
+ If the primary source of spin-up was orbital. we would not expect a strong correlation with the pulsed Lux and we would expect to see both increasesand decreases in pulse frequency as the pulsar moved relative to the observer.," If the primary source of spin-up was orbital, we would not expect a strong correlation with the pulsed flux and we would expect to see both increasesand decreases in pulse frequency as the pulsar moved relative to the observer."
+The absorption of light by dust in galaxies presents à major challenge in measuring accurate star formation rates (SET) in galaxies.,The absorption of light by dust in galaxies presents a major challenge in measuring accurate star formation rates (SFRs) in galaxies.
+ Understanding star formation is essential to our understanding of the evolutionary history of the Universe on both local and global scales., Understanding star formation is essential to our understanding of the evolutionary history of the Universe on both local and global scales.
+ SERs are commonly derived from Ilo and ultraviolet (UV) luminosities. both of which are significantly allected by dust) obscuration.," SFRs are commonly derived from $\alpha$ and ultraviolet (UV) luminosities, both of which are significantly affected by dust obscuration."
+ Lhe basic approach to correcting for this relics on having an observed constraint on the level of the attenuation. such as the Balmer decrement (Seaton1979:Cardellietal.1989: or the -parameter (Calzetti.Ikinney&Storchi-BergmannSteidel2000) together with a model for the dependence of the attenuation (Seaton1979:CardelliFisehera&Dopita 2005).," The basic approach to correcting for this relies on having an observed constraint on the level of the attenuation, such as the Balmer decrement \citep{Stn:79, Cdl:89, Cal:01} or the $\beta$ -parameter \citep{Cal:94,MHC:99, AS:00} together with a model for the wavelength-dependence of the attenuation \citep{Stn:79,Cdl:89,Cal:01, Pir:04, Tuf:04, FD:05}."
+. Dust has the dual elfect of not only reducing the overall Iuminosity observed. from a galaxy but. also. imprinting wavelength dependant spectral variations (C'alzetti1997a)., Dust has the dual effect of not only reducing the overall luminosity observed from a galaxy but also imprinting wavelength dependant spectral variations \citep{Cal:97a}.
+. Dust absorbs short-waveleneth radiation and. re-emits this energy at longer wavelengths., Dust absorbs short-wavelength radiation and re-emits this energy at longer wavelengths.
+ Dust also plavs à significant role in the fluid cvnamics. chemistry and. star. formation of galaxies (Drainectal.2007).," Dust also plays a significant role in the fluid dynamics, chemistry and star formation of galaxies \citep{Drn:07}."
+. Even though dust. plays a crucial role in galaxies. quantifving its effects and physical o»operties remains challenging.," Even though dust plays a crucial role in galaxies, quantifying its effects and physical properties remains challenging."
+ Understanding the roles jMaved: by inclination angles. optical depths and. galaxy morphologics when coupled with dust has been even more dillieult. and only recently have these effects been quantified ov Tulfsetal.(2004).," Understanding the roles played by inclination angles, optical depths and galaxy morphologies when coupled with dust has been even more difficult, and only recently have these effects been quantified by \citet{Tuf:04}."
+. Correcting for the elfects of dust attenuation is usually done using semi-empirical models (seeCarcellietal.1989:Byunοἱal.1994:Calzetti2001). but there are radiative ransfer (ICE) models developed for attenuation corrections (Nilourisetal.LOOT.1998.1999:Witt19Tulfs.al. 2004).," Correcting for the effects of dust attenuation is usually done using semi-empirical models \citep[see][]{Cdl:89,Byn:94,Cal:01} but there are radiative transfer (RT) models developed for attenuation corrections \citep{Xil:97,Xil:98,Xil:99,Wit:92,Tuf:04}."
+. Even though modelling techniques can give more accurate attenuation corrections than semi-empirical models. they require more observational information about the galaxies in order to produce accurate corrections on an object-by-object basis.," Even though modelling techniques can give more accurate attenuation corrections than semi-empirical models, they require more observational information about the galaxies in order to produce accurate corrections on an object-by-object basis."
+ However. the ICE techniques have proven extremely powerful when utilised to create libraries of model simulations CEullsetal.2004) for use in statistical analvsis of large statistical samples (Driveretal.2009).," However, the RT techniques have proven extremely powerful when utilised to create libraries of model simulations \citep{Tuf:04} for use in statistical analysis of large statistical samples \citep{Drv:07}."
+. An important. method for providing accurate SETts is he spectral energy distribution (SED) modelling technique., An important method for providing accurate SFRs is the spectral energy distribution (SED) modelling technique.
+ There are two main types of SED techniques: template itting SEDs (Calzettial.2000:Salimet2007) and ICE SED techniques (Silvaetal.1998:Popescu2000).," There are two main types of SED techniques: template fitting SEDs \citep{Cal:00,Slm:07} and RT SED techniques \citep{Slv:98, Pop:00}."
+. In xwticular the accuracy. of the ICE SED techniques is high. rut it tends to be complex and it requires knowledge of the SED over the whole electromagnetic spectrum. from UV to zu-infrared (PUR) and sub-millimeter wavelengths. which is not always available.," In particular the accuracy of the RT SED techniques is high, but it tends to be complex and it requires knowledge of the SED over the whole electromagnetic spectrum, from UV to far-infrared (FIR) and sub-millimeter wavelengths, which is not always available."
+ Due to this complexity the use of ICE echniques are less favoured than empirical methods for dust Correction., Due to this complexity the use of RT techniques are less favoured than empirical methods for dust correction.
+ Balmer decrements are commonly used as an indicator of obscuration corrections and are usually applied in conjunction with an extinction. law. such as those for the Milkv Way (MW) (Nandyetal.1975:Seaton1979:Cardellietal. 1989).. Small Magellanic Cloud. (τόνοιetal.1984:Bouchetet1955. SMC).. Large. Magellanic cloud (Fitzpatricketal.1986.LMC).. that. observed. for A31 (Bianchietal.1996).. starburst) galaxies. 1997a.2001) and other galaxies (Fischera&Dopita2005.ED05hereafter)...," Balmer decrements are commonly used as an indicator of obscuration corrections and are usually applied in conjunction with an extinction law such as those for the Milky Way (MW) \citep{Nnd:75,Stn:79, Cdl:89}, Small Magellanic Cloud \citep[SMC]{Prv:84, Bch:85}, Large Magellanic cloud \citep[LMC]{Ftz:86}, that observed for M31 \citep{Bnc:96}, starburst galaxies \citep{Cal:97a,Cal:01} and other galaxies \citep[FD05 hereafter]{FD:05}."
+ Any extinction law can be used to correct either nebular emission lines or continuum emission., Any extinction law can be used to correct either nebular emission lines or continuum emission.
+ As a consequence of how these curves are derived. though. some curves are preferred. over others to correct. for each type of emission. as the processes that result in these two types of emission are dillerent and it has been shown that nebular lines are usually obscurecd more than the continuum emission (Fanellietal.1988:Calzetti.Ixinney&Storchi-Dergmann1994:C'alzetti1997b:Alas-Llesse&Ixunth 1999).," As a consequence of how these curves are derived, though, some curves are preferred over others to correct for each type of emission, as the processes that result in these two types of emission are different and it has been shown that nebular lines are usually obscured more than the continuum emission \citep{Fan:88, Cal:94, Cal:97b, MK:99}."
+ A physical interpretation for this is provided by Keel(1993) and Calzetti.Ixinnevy&Storchi-Beremann(1094) who argue that the ionizing hot (voung) stars which produce the nebular lines are found in or close to the custy molecular clouds from which they were born. while the UV continuum is à product of older stars that have over time moved away from the dust clouds in which they formed. or have destroved the dust in situ revealing the cluster of stars.," A physical interpretation for this is provided by \citet{Kel:93} and \citet{Cal:94}
+ who argue that the ionizing hot (young) stars which produce the nebular lines are found in or close to the dusty molecular clouds from which they were born, while the UV continuum is a product of older stars that have over time moved away from the dust clouds in which they formed, or have destroyed the dust in situ revealing the cluster of stars."
+ The obscuration curves are also. dependent. on their environments. in particular the metallicity of the host galaxy (Cardellietal.1989).," The obscuration curves are also dependent on their environments, in particular the metallicity of the host galaxy \citep{Cdl:89}."
+. For instance. the LAIC. SAIC and MW curves have a large variation in values at the PUY end of the spectrum due to the dillerent metallicities of the three ealaxics and grain size distribution (C'alzetti.lxinnev&Storchi-Dergmann1994:Calzetti1997b). in the interstellar medium.," For instance, the LMC, SMC and MW curves have a large variation in values at the FUV end of the spectrum due to the different metallicities of the three galaxies and grain size distribution \citep{Cal:94,Cal:97b} in the interstellar medium."
+ These LMC€. SAIC ancl MW. obseuration. curves only take into account the dust between the observer and the star but in the case of external galaxies the dust. geometry is more complicated.," These LMC, SMC and MW obscuration curves only take into account the dust between the observer and the star but in the case of external galaxies the dust geometry is more complicated."
+ Not only does the cust between the observer and the stars of the galaxy need to be taken into account. but. back-scattering of light into the line of sight [rom dust in other regions of the galaxy must. also be considered and the dust is never uniformly. distributed.," Not only does the dust between the observer and the stars of the galaxy need to be taken into account, but back-scattering of light into the line of sight from dust in other regions of the galaxy must also be considered and the dust is never uniformly distributed."
+ The Calzetti extinction curve attempts to consider all these issues by folding a variety of obscuration ellects into a single expression., The Calzetti extinction curve attempts to consider all these issues by folding a variety of obscuration effects into a single expression.
+ Aleurer.IHleckmanandCalzetti(1999) show a relationship between the ratio of far-L (ELI) and UV Iuxes and the UV spectral slope 3 for a sample of starburst galaxies., \citet{MHC:99} show a relationship between the ratio of far-IR (FIR) and UV fluxes and the UV spectral slope $\beta$ for a sample of starburst galaxies.
+ Because this technique relates the FUR and UV radiation emitted from galaxies it can be a powerful tool in recovering the UV radiation lost due to the dust. regardless of the geometry of the dust.," Because this technique relates the FIR and UV radiation emitted from galaxies it can be a powerful tool in recovering the UV radiation lost due to the dust, regardless of the geometry of the dust."
+ The 3 parameter is based on the relation between the gradient in the UV. obscuration curve and the UV wavelengths. (Meurer..Heckman.andAdelberger&Steidel 2000).," The $\beta$ parameter is based on the relation between the gradient in the UV obscuration curve and the UV wavelengths \citep{MHC:99, Cal:94, AS:00}."
+.. However. this method is only recommended: for starburst galaxies as quiescent galaxies tend to deviate from the total FER to UV. luminosity ratio and UV spectral slope relation seen for starbursts (long 2004).," However, this method is only recommended for starburst galaxies as quiescent galaxies tend to deviate from the total FIR to UV luminosity ratio and UV spectral slope relation seen for starbursts \citep{Kng:04}."
+. The accuracy of this method is highly. dependent on the method emploved to measure the UV spectral slope., The accuracy of this method is highly dependent on the method employed to measure the UV spectral slope.
+ Calzetti.Winney&Storchi-Dergmann(1994) used LO bancs along the observed frame UV. continuum spectrum avoiding all [large scale features that deviate from the trend of the slope.," \citet{Cal:94}
+ used 10 bands along the observed frame UV continuum spectrum avoiding all large scale features that deviate from the trend of the slope."
+ With large scale surveys such as GALEN this level of accuracy is unattainable. leading to less accurate measurements of the slope and potentially Hawed: dust corrections.," With large scale surveys such as GALEX this level of accuracy is unattainable, leading to less accurate measurements of the slope and potentially flawed dust corrections."
+ Once suitable dust. corrections are in. hand. we can, Once suitable dust corrections are in hand we can
+the transfer would be conservative as usual iut large scale inass loss would be the domiuaut cause ol the A-terin in the ephemeris.,the transfer would be conservative as usual but large scale mass loss would be the dominant cause of the $A$ -term in the ephemeris.
+ Hypothesis (d) would be the most economical iuterpretatio1 but conflicts with the impersonal svutlieses of the light Clrves., Hypothesis (d) would be the most economical interpretation but conflicts with the impersonal syntheses of the light curves.
+ At this time. we favor a combination of interpretations (a) and (bJ.," At this time, we favor a combination of interpretations (a) and (b)."
+ In suiunary. our ststudy of the orbital »eriod ald the light curves reveals that CL Aur In a ‘Lasscal Aleol-tvpe interacting system witl the less massive and cool secondary star filling its inner Roele lobe.," In summary, our study of the orbital period and the light curves reveals that CL Aur is a classical Algol-type interacting system with the less massive and cool secondary star filling its inner Roche lobe."
+ ΤΙe possibility of a liot-spot mocle due to impact of streamline eas onto tle hot ar as lec| to confusiue cifficulties that we Iave resoved only tentatively., The possibility of a hot-spot model due to impact of streaming gas onto the hot star has led to confusing difficulties that we have resolved only tentatively.
+ High-resolution. near-IR spectroscoow should ‘eveal the ines of the secondary star auc lead to accurate absolute parameters o replace our estimates.," High-resolution, near-IR spectroscopy should reveal the lines of the secondary star and lead to accurate absolute parameters to replace our estimates."
+ Moreover. there is aso the possibility of obtaining direct evideuce of activity. such as complex aud variale liue profiles of Ha or variatious in the streneth of ie OL 7771 absorption line.," Moreover, there is also the possibility of obtaining direct evidence of mass-transfer activity, such as complex and variable line profiles of $\alpha$ or variations in the strength of the O I 7774 absorption line."
+ The evolutiona'N status of the system will then be more conviuciugly in haud., The evolutionary status of the system will then be more convincingly in hand.
+ The uore general reality of things is that the‘e are many short-period Aleols which have uot veel stuied so deeply as tlis present work has examined CL Aur., The more general reality of things is that there are many short-period Algols which have not been studied so deeply as this present work has examined CL Aur.
+ It would make a significant advance if more of these - having cillerent algeMaie sigus for the secular term of the period variability - could be brouglt to the same level of knowledge as the present binary., It would make a significant advance if more of these - having different algebraic signs for the secular term of the period variability - could be brought to the same level of knowledge as the present binary.
+ For iusta should all such binaries with A> 0 require a spo ou the |οἱ star near tle line of ceutewd would be major support for the reasoning conceruiug Wass 1joveinens.," For instance, should all such binaries with $A >$ 0 require a spot on the hot star near the line of centers, there would be major support for the reasoning concerning mass movements."
+ Should it also hap»en t the spot luminosity was cousisteutly found to be ower tha ithe kitetic iiipact requirect energy conversion. there woul be good reason to believe tha the majority o{the mass lost rom cool secondaries is. lu fact. lost to the systems aud not transerred conservaively.," Should it also happen that the spot luminosity was consistently found to be lower than the kinetic impact required for energy conversion, there would be good reason to believe that the majority of the mass lost from the cool secondaries is, in fact, lost to the systems and not transferred conservatively."
+ The shor-peri binaries found to require oxe | (seimi-detached systems with the primar vostars filline is lnuer Roche lobes) for their represeitatious would be expected. tJen. 10 SHOW SLjall-scale mass rausler to the secoucaries aud systemic mass loss [roin the prima‘ies.," The short-period binaries found to require mode 4 (semi-detached systems with the primary stars filling its inner Roche lobes) for their representations would be expected, then, to show small-scale mass transfer to the secondaries and systemic mass loss from the primaries."
+ Much. valiable observatioial aud mocleling work remaius., Much valuable observational and modeling work remains.
+ We would like to hauk the stalf ofthe Chungbuk National University Observatory [or assistaice with our observations. and Mr. Jae-Rim Isoo for the spectroscopic observations of CL Aur.," We would like to thank the staff of the Chungbuk National University Observatory for assistance with our observations, and Mr. Jae-Rim Koo for the spectroscopic observations of CL Aur."
+ We appreciate the carefu readiug aud valuable comments of the anonymous referee., We appreciate the careful reading and valuable comments of the anonymous referee.
+We have used he Sunbad data base maintained at CDS many times and appreciate its availability.,We have used the Simbad data base maintained at CDS many times and appreciate its availability.
+ This work has been done as part of a cooperative project between Chuugbuk National University aud the Ixo‘Ca Astronomy ancl Space Science Institute., This work has been done as part of a cooperative project between Chungbuk National University and the Korea Astronomy and Space Science Institute.
+ C.-H. Ix. was supported by the Ixorea Research Founcdatjon (SRE) Grant funded by the Ixorea government (MEST) (No., C.-H. K. was supported by the Korea Research Foundation (KRF) Grant funded by the Korea government (MEST) (No.
+ 2009-0069330)., 2009-0069330).
+"MSRPs during some portion of their evolution lose energy through a dominant mechanism other than magnetic dipole radiation (e.g. multipole radiation, gravitational wave or neutrino emission), then the evolution of pulsars through the P— diagram could be complex.","MSRPs during some portion of their evolution lose energy through a dominant mechanism other than magnetic dipole radiation (e.g. multipole radiation, gravitational wave or neutrino emission), then the evolution of pulsars through the $P-\dot{P}$ diagram could be complex."
+ A combination of the above mentioned factors (i.e. alternative progenitors and subsequent radiation) are then likely to play a role in millisecond pulsar evolution., A combination of the above mentioned factors (i.e. alternative progenitors and subsequent non-standard radiation) are then likely to play a role in millisecond pulsar evolution.
+ A MSXP period distribution that has sharp multimodal features coupled with non-standard energy loss mechanisms may be able to reconcile for the joint P— distribution of millisecond pulsars., A MSXP period distribution that has sharp multimodal features coupled with non-standard energy loss mechanisms may be able to reconcile for the joint $P-\dot{P}$ distribution of millisecond pulsars.
+ 'The research presented here has made use of the August 2008 version of the ATNF Pulsar Catalogue (Manchesteretal.|2005).., The research presented here has made use of the August 2008 version of the ATNF Pulsar Catalogue \citep{MHT93}.
+ The authors acknowledge NSF grant AST-0506453., The authors acknowledge NSF grant AST-0506453.
+"initial matter distribution, which guarantees that the large-scale structures closely resemble our local neighbourhood. but they differ in their random component. yielding different constrained realizations.","initial matter distribution, which guarantees that the large-scale structures closely resemble our local neighbourhood, but they differ in their random component, yielding different constrained realizations."
+ This allows us to check the robustness of our results with respect to the specitic realization., This allows us to check the robustness of our results with respect to the specific realization.
+ Finally. the two models also have different underlying cosmological models WMAP 3-year vs. WMAP 5-year best fit).," Finally, the two models also have different underlying cosmological models (WMAP 3-year vs. WMAP 5-year best fit)."
+ However. the effect of the background cosmology is well understood and results in an overall shift of the reionization earlier or later. with no significant effects on our results. for which only the relative timing of structure formation vs. reionization history is of importance.," However, the effect of the background cosmology is well understood and results in an overall shift of the reionization earlier or later, with no significant effects on our results, for which only the relative timing of structure formation vs. reionization history is of importance."
+ Our results show that the assumed efficiency of the ionizing sources has the most important influence on the nature of the reionization history of our Local Group of galaxies., Our results show that the assumed efficiency of the ionizing sources has the most important influence on the nature of the reionization history of our Local Group of galaxies.
+ Efficient photon production ensures that the nearby clusters emit more than sufficient number to ionize both themselves and their surroundings. including the Local Group.," Efficient photon production ensures that the nearby clusters emit more than sufficient number to ionize both themselves and their surroundings, including the Local Group."
+ The fact that those galaxy clusters (Virgo and Fornax) coincide with high. rare peaks of the density field means that they form their progenitor halos earlier than the LG. which is in a more average region of the universe.," The fact that those galaxy clusters (Virgo and Fornax) coincide with high, rare peaks of the density field means that they form their progenitor halos earlier than the LG, which is in a more average region of the universe."
+ As a result. the large-scale ionization fronts which propagated outward from the proto-clusters overrun the Local Group before it managed to form enough sources to ionize itself. resulting in its reionization being mostly externally-driven.," As a result, the large-scale ionization fronts which propagated outward from the proto-clusters overrun the Local Group before it managed to form enough sources to ionize itself, resulting in its reionization being mostly externally-driven."
+ Several points are worth noting here., Several points are worth noting here.
+ Although generally the radiative transfer is a highly non-local phenomenon. which feature complicates its numerical treatment and the code parallelization. during most of the EoR the situation. is. somewhat more complicated.," Although generally the radiative transfer is a highly non-local phenomenon, which feature complicates its numerical treatment and the code parallelization, during most of the EoR the situation is somewhat more complicated."
+ The neutral patches have enormous optical depth to soft ionizing radiation (the only type of radiation we consider here)., The neutral patches have enormous optical depth to soft ionizing radiation (the only type of radiation we consider here).
+ Even the already-ionized patches still have considerable continuum optical depth over cosmological (multiple Mpc) distances due to the small residual neutral fraction still remaining in suchregions., Even the already-ionized patches still have considerable continuum optical depth over cosmological (multiple Mpc) distances due to the small residual neutral fraction still remaining in such.
+ This residual neutral fraction diminishes over time. but does so only gradually. as more and more sources appear and the mean flux thereby increases.," This residual neutral fraction diminishes over time, but does so only gradually, as more and more sources appear and the mean flux thereby increases."
+ As a consequence of all this. reionization starts out as a fairly local process where only the relatively nearby. directly visible ionizing sources within the same ionized bubble contribute to the flux at a given point.," As a consequence of all this, reionization starts out as a fairly local process where only the relatively nearby, directly visible ionizing sources within the same ionized bubble contribute to the flux at a given point."
+ This property allows us to focus our analysis on the important local sources and ignore the far-away ones for our current purposes (they are of course all included in the radiative transfer simulation)., This property allows us to focus our analysis on the important local sources and ignore the far-away ones for our current purposes (they are of course all included in the radiative transfer simulation).
+ In our ~100 Mpe box there are multiple proto-clusters which collapse nonlinearly by the present. but of those only Virgo and Fornax are sufficiently close to potentially contribute to the reionization of our Local Group.," In our $\sim100$ Mpc box there are multiple proto-clusters which collapse nonlinearly by the present, but of those only Virgo and Fornax are sufficiently close to potentially contribute to the reionization of our Local Group."
+ Furthermore. the ionization fronts propagate through underdense regions (voids) much faster than. through. overdense ones (filaments. knots).," Furthermore, the ionization fronts propagate through underdense regions (voids) much faster than through overdense ones (filaments, knots)."
+ Therefore. the relative positioning of the structures of interest and the density fluctuations in their immediate neighbourhood are important.," Therefore, the relative positioning of the structures of interest and the density fluctuations in their immediate neighbourhood are important."
+ Once the available observational constraints are imposed in order to reproduce the local structures. we find that the Local Group is separated from Virgo and Fornax by voids in either realization (see Figure 29).," Once the available observational constraints are imposed in order to reproduce the local structures, we find that the Local Group is separated from Virgo and Fornax by voids in either realization (see Figure \ref{dens_slice_fig}) )."
+ In contrast. the previous studies of this problem which did not use constrained realizations (22) sampled a wide range of environments and relative positions of nearby clusters.," In contrast, the previous studies of this problem which did not use constrained realizations \citep{2007MNRAS.381..367W,2009ApJ...703L.167A}
+ sampled a wide range of environments and relative positions of nearby clusters."
+ Such. purely statistical approach yields valuable insights on the range of reionization histories that could be expected for a certain type object (e.g. LG-like objects).," Such, purely statistical approach yields valuable insights on the range of reionization histories that could be expected for a certain type object (e.g. LG-like objects)."
+ However. by its nature such approach necessarily includes many objects which. although they share certain basic features. locally do not reproduce the specitic large-scale structures around us.," However, by its nature such approach necessarily includes many objects which, although they share certain basic features, locally do not reproduce the specific large-scale structures around us."
+ Therefore. the constrained realizations are indispensible if we want to make realiable predictions for the effects of reionization on our neighbourhood.," Therefore, the constrained realizations are indispensible if we want to make realiable predictions for the effects of reionization on our neighbourhood."
+ Why is the mode of reionization. external vs. internal. of our Local Group an important issue?," Why is the mode of reionization, external vs. internal, of our Local Group an important issue?"
+ This has a number of important implications for the formation of structures., This has a number of important implications for the formation of structures.
+ Reionization dramatically rises the Jeans mass. thus impeding the formation and growth of small galaxies.," Reionization dramatically rises the Jeans mass, thus impeding the formation and growth of small galaxies."
+ In terms of this effect. the galactic haloes fall into three categories.," In terms of this effect, the galactic haloes fall into three categories."
+ The gas in the smallest halos (ninihalos). whose virial temperatures are below the limit (~10! K) for efficient radiative cooling through atomic line radiation.," The gas in the smallest halos (minihalos), whose virial temperatures are below the limit $\sim10^4$ K) for efficient radiative cooling through atomic line radiation."
+ The ionization of the gas brings its temperature to ~107 K and it boils out. resulting in their complete evaporation €?2).. which leaves behind dark halos.," The ionization of the gas brings its temperature to $\sim10^4$ K and it boils out, resulting in their complete evaporation \citep{2004MNRAS.348..753S,2005MNRAS...361..405I}, which leaves behind dark halos."
+ In the other limit. the galaxies above certain mass (AL>1017AL. ) have sufficiently deep gravitational potential wells to sucessfully withstand the effects of ionizing radiation and are thus not significantly affected by the reionization process.," In the other limit, the galaxies above certain mass $M\gtrsim10^{10}M_\odot$ ) have sufficiently deep gravitational potential wells to sucessfully withstand the effects of ionizing radiation and are thus not significantly affected by the reionization process."
+ The effects of radiative feedback on dwarf galaxies of intermediate mass. roughly between 10737. and 102AZ. is more complex and still very much a subject of active investigation.," The effects of radiative feedback on dwarf galaxies of intermediate mass, roughly between $10^8M_\odot$ and $10^{10}M_\odot$ is more complex and still very much a subject of active investigation."
+ The gas in such already-formed systems cannot be photoevaporated. as it can cool back down to ~10! K very efficiently.," The gas in such already-formed systems cannot be photoevaporated, as it can cool back down to $\sim10^4$ K very efficiently."
+ However. jfhetoionization heating rises the intergalactic gas temperature and oessure. which rises the Jeans mass and thereby suppresses the uture formation of very low-mass galaxies. as well as curtails he fresh gas infall onto such halos.," However, photoionization heating rises the intergalactic gas temperature and pressure, which rises the Jeans mass and thereby suppresses the future formation of very low-mass galaxies, as well as curtails the fresh gas infall onto such halos."
+ Larger galaxies are less affected directly. but could do so indirectly. through their smaller wrogenitors. which could be expected e.g. to result in smoother gas sub-structure and moditied stellar populations.," Larger galaxies are less affected directly, but could do so indirectly, through their smaller progenitors, which could be expected e.g. to result in smoother gas sub-structure and modified stellar populations."
+ Where the boundary between efficient and inefficient feedback from reionization lies is still unclear and very much subject of active research., Where the boundary between efficient and inefficient feedback from reionization lies is still unclear and very much subject of active research.
+ Ful investigation of the effects of reionization on galaxy formation and satellite galaxy populations goes well beyond the scope of the current work., Full investigation of the effects of reionization on galaxy formation and satellite galaxy populations goes well beyond the scope of the current work.
+ However. our present results indicate that the photon production efficiencies of the first galaxies are the main factor determining the type of reionization history which our Loca Group underwent.," However, our present results indicate that the photon production efficiencies of the first galaxies are the main factor determining the type of reionization history which our Local Group underwent."
+ Therefore. this process should have left usefu fossil records in the properties of our neighbourhood which wil help us use local observations to answer some of the Key questions about the young universe.," Therefore, this process should have left useful fossil records in the properties of our neighbourhood which will help us use local observations to answer some of the key questions about the young universe."
+ This study was supported in part by Swiss National Science Foundation grant 200021-116696/1 and Swedish Research Council grant 60336701., This study was supported in part by Swiss National Science Foundation grant 200021-116696/1 and Swedish Research Council grant 60336701.
+ GY acknowledges support of MICINN (Spain) through research. grants PPA2009-O8958. AYA2009-|3875-C03-02 and CONSOLIDER-INGENIO SyEC ¢CSD2007.0050).," GY acknowledges support of MICINN (Spain) through research grants FPA2009-08958, AYA2009-13875-C03-02 and CONSOLIDER-INGENIO SyEC (CSD2007.0050)."
+ YH. has been partially supported by the ISF (13/08)., Y.H. has been partially supported by the ISF (13/08).
+ The CLUES simulations have been performed in the MareNostrum supercomputer at BSC (Spain) and in the HLRBII Altix computer at LRZ (Germany)., The CLUES simulations have been performed in the MareNostrum supercomputer at BSC (Spain) and in the HLRBII Altix computer at LRZ (Germany).
+ We also thank DEISA for granting us epu time in these computers through two DECT projects SIMU-LU and SIMUGAL-LU., We also thank DEISA for granting us cpu time in these computers through two DECI projects SIMU-LU and SIMUGAL-LU.
+ We thank Nick Gnedin for making publicly available his visualization code IFRIT. which was used to produce the images in Figs.," We thank Nick Gnedin for making publicly available his visualization code IFRIT, which was used to produce the images in Figs."
+ +6., 4-6.
+ Some of the radiative transfer simulations, Some of the radiative transfer simulations
+line in regions of intermediate or low optical depth.,time in regions of intermediate or low optical depth.
+" In section 5.5 it is shown that a propagating ""adiatiou front can be more accurately evolved in this situation when the cilfusion coellicients are recalculated every diffusion substep.", In section \ref{sec:front} it is shown that a propagating radiation front can be more accurately evolved in this situation when the diffusion coefficients are recalculated every diffusion substep.
+ Coisidering E7;1 as a siugle vector of leugth ΑκAL. the updated values may be determined w solviig an equation with a matix of CNxΑΙ] οenients. using a sparse-1uatrix technique such as LUCOLiplete Cholesky — conjugate eracien (Press e al.," Considering $E^{n+1}_{i,j}$ as a single vector of length $N\times M$, the updated values may be determined by solving an equation with a matrix of $(N\times M)^2$ elements, using a sparse-matrix technique such as incomplete Cholesky – conjugate gradient (Press et al."
+ 1992)., 1992).
+ These methods liave proven useful Or partieular problems with fixed Xucdary couditiois. but the form of the matrix and tlie uethod required vary with the boudary cor(itions.," These methods have proven useful for particular problems with fixed boundary conditions, but the form of the matrix and the sparse-matrix method required vary with the boundary conditions."
+ Solutious of equation (25)) may also ye obtaiued by combining simple solvers wit1 multi-g4d acceleration techuiques (Hackbusch 1985)., Solutions of equation \ref{eqn:delfluxdifferencing}) ) may also be obtained by combining simple solvers with multi-grid acceleration techniques (Hackbusch 1985).
+ Howeve. here we choose an alternatiug-cli'ectiou-itlicit. (ADI) method on a single eril.," However, here we choose an alternating-direction-implicit (ADI) method on a single grid."
+ This uay be less efficient than the best sparse-jalrlix al[1 uulti-grid algorithius. though the uumber of operations required is proportional to NxAZ.," This may be less efficient than the best sparse-matrix and multi-grid algorithms, though the number of operations required is proportional to $N\times M$."
+" In ADL. successive approximatious to E"".1 are computed by advauciug towards a t-5tatioary state he equation The uew variable w may be thought of as the pseudo-time."," In ADI, successive approximations to $E^{n+1}$ are computed by advancing towards a $w$ -stationary state the equation The new variable $w$ may be thought of as the pseudo-time."
+ Each w-step is split into two parts., Each $w$ -step is split into two parts.
+ Iu the first. the update is w-imiplicit along the 1-directiou. and explicit aloug the 2-directiou.," In the first, the update is $w$ -implicit along the 1-direction, and explicit along the 2-direction."
+ Iu the second. the differeucing schemesfor tlie two axes are exchanged.," In the second, the differencing schemesfor the two axes are exchanged."
+ Labeling the approximate value ol E? at pseudo-timestep 11 by E75. the dillerence equations solved are ol sweeps which: are implicit aloug the 1-direction. aud ol sweeps which are implicit along the 2-direction.," Labeling the approximate value of $E^{n+1}_{i,j}$ at pseudo-timestep $m$ by $E^m_{i,j}$, the difference equations solved are on sweeps which are implicit along the 1-direction, and on sweeps which are implicit along the 2-direction."
+ Each sweep involves solving a tridiagonal matrix equation., Each sweep involves solving a tridiagonal matrix equation.
+ For example. for boundary coudition VE= 0. the matrix equation solved on the," For example, for boundary condition $\nabla E=0$ , the matrix equation solved on the"
+"Oct. 2004 show excess absorption (Gendreetal.,2006).",Oct. 2004 show excess absorption \citep{gcp06}.
+". In a systematic study of 93 promptly observedSwift GRBs with known redshift (up to May 2009), 85 show evidence of intrinsic X-ray absorption at the host galaxy site (Campanaetal.,2010)."," In a systematic study of 93 promptly observed GRBs with known redshift (up to May 2009), 85 show evidence of intrinsic X-ray absorption at the host galaxy site \citep{ctu10}."
+". Similarly, in our sample, we detect excess absorption in 26 out of 33 cases."," Similarly, in our sample, we detect excess absorption in 26 out of 33 cases."
+" This difference in the excess absorption detection rate is primarily related to the quality of the X-ray spectra, whichpre-Swift was typically taken 8—12 hours post trigger, by which time the signal-to-noise ratio was insufficiently high to accurately measure any intrinsic absorption."," This difference in the excess absorption detection rate is primarily related to the quality of the X-ray spectra, which was typically taken 8–12 hours post trigger, by which time the signal-to-noise ratio was insufficiently high to accurately measure any intrinsic absorption."
+" Figure 5 shows a comparison of the effective neutral hydrogen absorption Ny and visual extinction Ay, both in the rest frame of the corresponding burst."," Figure \ref{avnh} shows a comparison of the effective neutral hydrogen absorption $N_{\rm H}$ and visual extinction $A_{\rm V}$, both in the rest frame of the corresponding burst."
+" As has been noted frequently in the past (e.g.Galama&Wijers,2001;Strattaetal.,2004;Schady 2010),, the Ny-to-Ay ratio is far from being similar among different bursts, and substantially larger than in our Galaxy."," As has been noted frequently in the past \citep[e.g.][]{gaw01, sfa04, spo10}, the $N_{\rm H}$ $A_{\rm V}$ ratio is far from being similar among different bursts, and substantially larger than in our Galaxy."
+" Note, however, that here, as has been usual in previous cases, solar metallicity has been assumed in deriving Ny."," Note, however, that here, as has been usual in previous cases, solar metallicity has been assumed in deriving $N_{\rm H}$."
+" Since the observed curvature in the X-ray spectra is predominantly, but not exclusively, due to absorption by oxygen, the derived effective Ny is inversely proportional to the metallicity (or better O/H ratio) of the burst environment."," Since the observed curvature in the X-ray spectra is predominantly, but not exclusively, due to absorption by oxygen, the derived effective $N_{\rm H}$ is inversely proportional to the metallicity (or better O/H ratio) of the burst environment."
+" Since this metallicity has been observed (in other GRBs) to be about 1/5 solar (though extremes of solar (Prochaska up to super-solar al.,2010) and nearly 1/100 solar (D’Eliaetal.,2007;Raual,2010) do also occur), the effective Ny would likely be even larger than shown in Fig. 5,,"," Since this metallicity has been observed (in other GRBs) to be about 1/5 solar (though extremes of solar \citep{psp09,efh09} up to super-solar \citep{srg10} and nearly 1/100 solar \citep{dfm07, rsk10} do also occur), the effective $N_{\rm H}$ would likely be even larger than shown in Fig. \ref{avnh},"
+ if the proper line-of-sight metallicity were to be used (if it were known)., if the proper line-of-sight metallicity were to be used (if it were known).
+" In contrastto our Ay distribution, the distribution of Nux from the completeSwift sample lacks a substantial fraction of zero column density (Campanaetal,2010)."," In contrastto our $A_{\rm V}$ distribution, the distribution of $N_{\rm H,X}$ from the complete sample lacks a substantial fraction of zero column density \citep{ctu10}."
+". This has been explained by Campanaetal.(2010) as evidence that the bursts originate within high-density regions of their hosts, since a random distribution in a galaxy like ours would predict a sizable fraction (~30%)) with no intrinsic absorption."," This has been explained by \cite{ctu10} as evidence that the bursts originate within high-density regions of their hosts, since a random distribution in a galaxy like ours would predict a sizable fraction $\sim$ ) with no intrinsic absorption."
+" By combining their sample with the Lyman-a absorbers at z>2 of Fynboetal.(2009b) they also find, similar to earlier reports (e.g.Watsonetal.,2007),, that the bulk of GRBs have column densities in X-rays which are a factor ~10 higher than in the optical (Nur), which they explain by ionization of hydrogen by the high energy flux of the GRB."," By combining their sample with the $\alpha$ absorbers at $z>2$ of \cite{fjp09} they also find, similar to earlier reports \citep[e.g.][]{whf07}, that the bulk of GRBs have column densities in X-rays which are a factor $\sim$ 10 higher than in the optical $N_{\rm HI}$ ), which they explain by ionization of hydrogen by the high energy flux of the GRB."
+" Since this ratio is roughly similar to that of Nyx vs. Ay (Fig. 5)),"," Since this ratio is roughly similar to that of $N_{\rm H,X}$ vs. $A_V$ (Fig. \ref{avnh}) ),"
+ one could think that Ay and the Lyman-a absorption are correlated., one could think that $A_V$ and the $\alpha$ absorption are correlated.
+" However, the three bursts with reported Ny; in Fynboetal.(2009b) (070802, 071031 and 080804) do not show any correlation, similar to the 6 bursts from the UVOT sample published by Schadyetal. (2010)."," However, the three bursts with reported $N_{\rm HI}$ in \cite{fjp09} (070802, 071031 and 080804) do not show any correlation, similar to the 6 bursts from the UVOT sample published by \cite{spo10}."
+. The redshift distribution of our sample is complete (also one of the 4 GRBs not detected by GROND has a redshift)., The redshift distribution of our sample is complete (also one of the 4 GRBs not detected by GROND has a redshift).
+" We re-iterate that the only selection criterion was the detection of an X-ray afterglow, and do not see a bias introduced by the requirement of a rapid GROND observation (or equivalently an occurrence during Chilean night time)."," We re-iterate that the only selection criterion was the detection of an X-ray afterglow, and do not see a bias introduced by the requirement of a rapid GROND observation (or equivalently an occurrence during Chilean night time)."
+ A comparison to the distribution of all known long-duration bursts (about complete; Fig. 6)), A comparison to the distribution of all known long-duration bursts (about complete; Fig. \ref{zdis}) )
+" reveals the former to have a flatter distribution, with a somewhat higher number of z>4 bursts."," reveals the former to have a flatter distribution, with a somewhat higher number of $z>4$ bursts."
+" However, a KS-test shows that this is not statistically significant, and both distributions are consistent with being drawn from the same sample within lo."," However, a KS-test shows that this is not statistically significant, and both distributions are consistent with being drawn from the same sample within $\sigma$."
+" The presence of GRBs 080913 and 090429B in our sample corresponds to a fraction of 5.542.896 of bursts at redshifts 5, and the strict upper limit would be if all three GRBs without redshift would be at z>5."," The presence of GRBs 080913 and 090429B in our sample corresponds to a fraction of $\pm$ of bursts at redshifts $z>5$ , and the strict upper limit would be if all three GRBs without redshift would be at $z>5$."
+ A larger sample size with a similarly good completeness level would be required to derive fractions with errors less than the present ~50% level., A larger sample size with a similarly good completeness level would be required to derive fractions with errors less than the present $\sim$ level.
+" The majority of afterglow SEDs show a spectral break between the X-ray and optical/NIR range that can be described with a slope difference of 0.5, consistent with the basic fireball scenario."," The majority of afterglow SEDs show a spectral break between the X-ray and optical/NIR range that can be described with a slope difference of 0.5, consistent with the basic fireball scenario."
+ This spectral break implies a R-band flux of about 2-- mag fainter than obtained from an extrapolation of the X-ray spectrum., This spectral break implies a $R$ -band flux of about 3--4 mag fainter than obtained from an extrapolation of the X-ray spectrum.
+" This effect is dealt with in all the recent definitions of “darkness” (e.g.Pedersenetal.,2006),, and our finding of a dominance of SED breaks is consistent with the large number of bursts seen at fo< 1."," This effect is dealt with in all the recent definitions of “darkness” \citep[e.g.][]{phh06}, and our finding of a dominance of SED breaks is consistent with the large number of bursts seen at $\beta_{\rm O} < 1$ ."
+" The faint optical afterglow emission of “dark bursts”, where we used the definition of vanderHorstetal.(2009),, is due to a mixture of moderate intrinsic extinction at moderate redshifts, and a fraction of bursts at redshift >5 (about of the dark bursts)."," The faint optical afterglow emission of “dark bursts”, where we used the definition of \cite{hkg09}, is due to a mixture of moderate intrinsic extinction at moderate redshifts, and a fraction of bursts at redshift $>$ 5 (about of the dark bursts)."
+" This finding is in line with previous investigations (Melandrietal.,2008;Cenko2009b;Perley2009; 2009b).."," This finding is in line with previous investigations \citep{mmk08, ckh09, pcb09, fjp09}. ."
+" In particular, Cenkoetal.(2009b) used a similar approachas ours, namely a sample of 29 bursts for which follow-up observations with the robotic Palomar"," In particular, \cite{ckh09} used a similar approachas ours, namely a sample of 29 bursts for which follow-up observations with the robotic Palomar"
+The ceutral DIT resides in a deuse and thick gaseous disk (Fig. 1)).,"The central BH resides in a dense and thick gaseous disk (Fig. \ref{fig1}) ),"
+ which could obscure an ACN., which could obscure an AGN.
+ To quantity lis. we computed the column densities over 60 random. ines of sight through the centers of galaxy IRI at #=330 Myr and WR2 at t=380 Myr.," To quantify this, we computed the column densities over 60 random lines of sight through the centers of galaxy HR1 at $t$ =330 Myr and HR2 at $t$ =380 Myr."
+ The results are insensitive to the choice of ceuter. selected to be either he ceuter of mass of old stars or the point of maximal vorticity.," The results are insensitive to the choice of center, selected to be either the center of mass of old stars or the point of maximum vorticity."
+ The foreground hydrogen column deusitv Vy was ostimated over a cross-section corresponding to the iiminmn Jeans lenueth. Ενeave.," The foreground hydrogen column density $N_\mathrm{H}$ was estimated over a cross-section corresponding to the minimum Jeans length, $4 \times \epsilon_\mathrm{AMR}$."
+ Figure Lo shows the distribution of Ny for runs IIRI and ITR2 and for various sets of inclinations., Figure \ref{fig4} shows the distribution of $N_\mathrm{H}$ for runs HR1 and HR2 and for various sets of inclinations.
+ of the ines of sight reach Compton thickness. loe(Nyy)c21.1. and have logCNqg)=23. Le. stroug obscuration.," of the lines of sight reach Compton thickness, $\log(N_\mathrm{H}) \geq 24.1$, and have $\log(N_\mathrm{H}) \geq 23$, i.e., strong obscuration."
+ The obscuration teuds to be higher iu the more massive disk IIR1. where Compton thickness can be reached even for face-on oricutations owiug to the ~1 kpe thickness of the eas disk.," The obscuration tends to be higher in the more massive disk HR1, where Compton thickness can be reached even for face-on orientations owing to the $\sim 1$ kpc thickness of the gas disk."
+ The lower-1uass galaxy IIR2. perhaps more representativo at 2«2. can reach Compton thickuess imost exclusively for edec-ou orieutatious.," The lower-mass galaxy HR2, perhaps more representative at $z<2$, can reach Compton thickness almost exclusively for edge-on orientations."
+ If most eas in the central parsec lies m a torus. uuresolved in our models. the obscuration ou edge-onu projections may be further enliauced aud the dependence of obscuration on inclination may be strougcr.," If most gas in the central parsec lies in a torus, unresolved in our models, the obscuration on edge-on projections may be further enhanced and the dependence of obscuration on inclination may be stronger."
+ Siguificaut obscuration would still be present alone most lines of sieht. because the coluun clensitics measured in the stations are on average dominated by the ceutral LO pe.," Significant obscuration would still be present along most lines of sight, because the column densities measured in the simulations are on average dominated by the central 10 pc."
+ The isolated simulations preseuted here. using high eas fraction representative of +~2 disks. have a disk inflow rate consistent with cosmological simulations of high-redshift stroaun-fed galaxies (Ceverinoetal.2010).," The isolated simulations presented here, using high gas fraction representative of $z\sim 2$ disks, have a disk inflow rate consistent with cosmological simulations of high-redshift stream-fed galaxies \citep{ceverino}."
+. They reach resolutions better than 2 pc. hence resolving the iuclear inflow down to scales at which other processes dive the actual simall-scale BIL accretion (Combes 200101.," They reach resolutions better than 2 pc, hence resolving the nuclear inflow down to scales at which other processes drive the actual small-scale BH accretion \citep{combes00}."
+.. Onulv a fraction of the inflowine gas needs to nake it all the way iuto the DIL., Only a fraction of the inflowing gas needs to make it all the way into the BH.
+ The issues are simular o the case of imerecr-driven fucling. namely ectting vid of angular momenta while avoiding excessive star ormation aud ACN feedback.," The issues are similar to the case of merger-driven fueling, namely getting rid of angular momentum while avoiding excessive star formation and AGN feedback."
+ This is bevoud the scope of our paper. and we limit ourselves to heuristic estimates.," This is beyond the scope of our paper, and we limit ourselves to heuristic estimates."
+ One can asstune that the local relation between the DII mass Mp aud the bulge properties (1299 Af. aud velocity dispersion 7} is crudely valid at high redshift., One can assume that the local relation between the BH mass $\MBH$ and the bulge properties (mass $\Mblg$ and velocity dispersion $\sigma$ ) is crudely valid at high redshift.
+ If we adopt MpMis~10ὁ at 2=0. assume it scales as σ aud allow a cosmological scaling a?x(11i. then we obtain AJpu/Mys~Bs10? at i~ 2," If we adopt $\MBH/\Mblg
+\sim 10^{-3}$ at $z\!=\!0$, assume it scales as $\sigma^2$ and allow a cosmological scaling $\sigma^2 \propto (1+z)$, then we obtain $\MBH/\Mblg \sim 3 \times 10^{-3}$ at $z\! \sim\! 2$."
+A ~~2 galaxy of barvonic mass 1005ME. hosts a DIE of ~l105ADS. while in a eravitationally unstable steady state it is typically half bulee aud half disk (DSC'09).," A $z\! \sim \! 2$ galaxy of baryonic mass $10^{11}\,\msun$ hosts a BH of $\sim \! 10^8\, \msun$, while in a gravitationally unstable steady state it is typically half bulge and half disk (DSC09)."
+" According to(2).. the average inflow rate through the disk into the bulee is Mj,~12.5Mvrο."," According to, the average inflow rate through the disk into the bulge is $\Mdotblg \sim 12.5 \sy$."
+ HE the ratio of average accretion rates into bulge aud BU follow the ratio of the correspondiug masses. the BIT accretes ou average at δη~O.OIALvrbt. which is of the Eddington rate (assuming a 0.1 cfiicicney for u1ass-to-cherey couversion).," If the ratio of average accretion rates into bulge and BH follow the ratio of the corresponding masses, the BH accretes on average at $\MdotBH \sim 0.04 \sy$, which is of the Eddington rate (assuming a $0.1$ efficiency for mass-to-energy conversion)."
+" The corresponding bolometric luuiuositv is 2«lüll'ergsο,", The corresponding bolometric luminosity is $2 \times 10^{44}\ergs$.
+" With typically in N-ravs.- we estimate: on average Lx~T107IqueyesL scaling: with: ealaxy mass and with: (1↽|25:)"".WhileAc the average huninosity would be modest. short episodes of higher accretion rate. possibly up to the Eddington level. occur during the central coalesccuce of mierating git chuups — which could also brine with them sced BUs (Ehucercenetal.2008a)."," With typically in X-rays, we estimate on average $L_{\rm X} \sim 10^{42-43}\ergs$, scaling with galaxy mass and with $(1+z)^{2.5}$.While the average luminosity would be modest, short episodes of higher accretion rate, possibly up to the Eddington level, occur during the central coalescence of migrating giant clumps – which could also bring with them seed BHs \citep{EBE08}."
+. At very high redshift. the siue process could feed brighter AGN in rare massive systems (DiMatteoctal. provided that unstable disks indeed form there.," At very high redshift, the same process could feed brighter AGN in rare massive systems \citep{DM11}, provided that unstable disks indeed form there."
+ The average inflow rate increases with redshift (eq., The average inflow rate increases with redshift (eq.
+ 2). such that a2 =610 disk can support coutiuuous accretion at the Exdiugtou rate in self-regulated quasar mode if Αιλα is ~O.OL0.08. for the same ~10? fraction of the inflow rate accreting onto the BIL and assununue outflows would not affect the cold gas streaming inward.," 2), such that a $z\!=\!6-10$ disk can support continuous accretion at the Eddington rate in self-regulated quasar mode if $\MBH/\Md$ is $\sim\!0.04-0.08$, for the same $\sim\!10^{-3}$ fraction of the inflow rate accreting onto the BH, and assuming outflows would not affect the cold gas streaming inward."
+ Then a seed DII o£ LAL. at z10 can erow oxponcutially by 10 c-folds to ὃς10AL. at 2~6. possibly explainius verv massive BUs erowing iu eas-rich systems at 22:5 (Mortlocketal.2011:Treisteral.2011).," Then a seed BH of $10^5\,\msun$ at $z \!\sim \! 10$ can grow exponentially by 10 e-folds to $2\times 10^9\,\msun$ at $z\! \sim\! 6$, possibly explaining very massive BHs growing in gas-rich systems at $z\!>\!5$ \citep{mortlock,treister11}."
+. The typical barvouic mass of the host ealaxy is ~1OMAL.. and if the halo mass is a few times 1017AL... the comoving umuber density of such bright +>5 ACNs would be ~LO“to 10.Mpe7," The typical baryonic mass of the host galaxy is $\sim \! 10^{11}\,\msun$, and if the halo mass is a few times $10^{12}\,\msun$, the comoving number density of such bright $z>5$ AGNs would be $\sim \!10^{-8}$ to $10^{-6}\,\Mpc^{-3}$."
+ Violent eravitational instability im high-redshift disk ealaxies naturally leads to the erowth of a bulge aud a central DILE it can both feed an ACN iud obscure it., Violent gravitational instability in high-redshift disk galaxies naturally leads to the growth of a bulge and a central BH; it can both feed an AGN and obscure it.
+ This stems from the developiug picture where many star-forming galaxies at high redshift are rotating disks with high eas fractions continuously fed by cold streams., This stems from the developing picture where many star-forming galaxies at high redshift are rotating disks with high gas fractions continuously fed by cold streams.
+ Such disks develop imstabilitv that involves transient perturbations aud giant clumps., Such disks develop instability that involves transient perturbations and giant clumps.
+ Cravitational torquiug amone these perturbations lead to mass inflow. which provides cucrey for maintaimine the strong turbulence required for selfreeulated instability at Q~1.," Gravitational torquing among these perturbations lead to mass inflow, which provides energy for maintaining the strong turbulence required for self-regulated instability at $Q\! \sim \! 1$."
+ The cosinological inflow rate along cold streams sets au upper luit. but the actual inflow rate toward the nucleus is determined by the disk iustabilitv ou ealactic scales.," The cosmological inflow rate along cold streams sets an upper limit, but the actual inflow rate toward the nucleus is determined by the disk instability on galactic scales."
+ Ileh-edshift disks are very clifferent from nearby spirals. the key difference being the auch ligher eas fraction.," High-redshift disks are very different from nearby spirals, the key difference being the much higher gas fraction."
+ The instability then operates on short dynamical timescales. aud drives an intense Ποπ; through the disk at the level of ~30MνιHAZ(1|Ja .," The instability then operates on short dynamical timescales, and drives an intense inflow through the disk at the level of $\sim \!10\sy M_{\rm b,11} (1+z)_3^{3/2}$ ."
+ Iu low-redshift disks. the instability is a secular process limited to non-axisviunetrce modes with considerably weaker torquiug: bars are mvoked iu some AGN models (Fanidakisetal.2011:Degehuau2006) but need about teu rotation periods to convey sole eas inward (Dournaudetal.2005).," In low-redshift disks, the instability is a secular process limited to non-axisymmetric modes with considerably weaker torquing: bars are invoked in some AGN models \citep{fanidakis, begelman} but need about ten rotation periods to convey some gas inward \citep{BCS05}."
+. Cravitational torquing in hiel-: disks is wach more efücent aud involves richer gas reservoirs., Gravitational torquing in $z$ disks is much more efficient and involves richer gas reservoirs.
+ Disk instability at +~2 can thus funnel half of the disk gas toward the center in 2 Cyr., Disk instability at $z \! \sim \! 2$ can thus funnel half of the disk gas toward the center in 2 Gyr.
+ This is similar to the mass inflow iu a 1najor mereer (IHopkiusctal.2006).. but spread over a ten times longer period. resulting iu a lower average ACN luminosity. with higher duty cvele. and lieh obscuration.," This is similar to the mass inflow in a major merger \citep{H06}, but spread over a ten times longer period, resulting in a lower average AGN luminosity, with higher duty cycle, and high obscuration."
+ This process could dominate. as the gas mass in the relatively sinooth cosmological streais is larger than that associated with mergers less than of the cosmological infall is iu major mergers— (Dekelctal.2000a:L'ITuilBeret 2011).," This process could dominate, as the gas mass in the relatively smooth cosmological streams is larger than that associated with mergers – less than of the cosmological infall is in major mergers \citep{dekel09, lhuillier11}."
+. The iain prediction is thus that many hieh-: ACNs should be hosted by star-forming disk ealaxies. composed of clampy," The main prediction is thus that many $z$ AGNs should be hosted by star-forming disk galaxies, composed of clumpy"
+qi.)=0.39. 0.02. 0.06. 0.10. 0.13. were found.,"$q(R,\varphi)=0.39,$ $0.02,$ $0.06,$ $-0.10,$ $-0.13$ were found."
+ Thus. Condition LL is indeed. satisfied.," Thus, Condition II is indeed satisfied."
+ The outlier value of 0.39 is produced by cancellation of Oel.=90.5ms7 and οz|28.2ms7. ie. the free fall of the upper atmosphere is for a short timealmost uniform with constant velocity and vanishing acceleration —g(A.ο)&2.3ms2," The outlier value of $0.39$ is produced by cancellation of $\partial v/\partial t\approx -30.5\:\mbox{ms}^{-2}$ and $v\partial v/\partial r\approx +28.2\:\mbox{ms}^{-2}$, i.e. the free fall of the upper atmosphere is for a short timealmost uniform with constant velocity and vanishing acceleration $-g(R,\varphi)\approx -2.3\:\mbox{ms}^{-2}$."
+ Vherefore. the density. stratification οους from. (7)) and his is manifested in the large q.," Therefore, the density stratification differs from \ref{2.110}) ) and this is manifested in the large $q$."
+ However.this phase could also be used to ect the best d and M4 in Table 20 because aqE<lms2," However,this phase could also be used to get the best $d$ and ${\cal M}$ in Table \ref{tab2} because $\vert a^{\rm (dyn)} \vert < 1\:\mbox{ms}^{-2}$."
+ In the interval 0.92ο< 0.95. the acceleration is dominated by go.z89.112.148.102ms7o»g. and he atmosphere is practically at a standstill.," In the interval $0.92 \leq \varphi \leq 0.95$ the acceleration is dominated by $g_{\rm e}\approx 89,112,148,102\:\mbox{ms}^{-2} \gg g_{\rm s}$ and the atmosphere is practically at a standstill."
+ Phe values q=1.06.L68.0.12.1.44 and aUPm(ndg.G.D)m5.28..2.6 −⋅⋠⋠indicate aid777z41.254.−19.55ms2 and (ilswyg..," The values $q=1.06,-1.68,-0.12,1.44$ and $a^{\rm (dyn)}(R,t)/g_{\rm s}(R,t)\approx 5,28,-2,6$ indicate $a^{\rm (dyn)}\approx 41,254,-19,55\:\mbox{ms}^{-2}$ and $\vert a^{\rm (dyn)} \vert \gg g_{\rm s}$."
+" Ug. is neglected. by setting Mid!)=0 and the data of Table 1. are used with a""?= 0. the upper limit d«642pc is obtained from y= 0.94."," If $g_{\rm s}$ is neglected by setting ${\cal M}(d,t)=0$ and the data of Table \ref{tab1} are used with $a^{\rm (dyn)}=0$ , the upper limit $d < 642\mbox{pc}$ is obtained from $\varphi=0.94$ ."
+ LW this estimation is improved. by aU--=094)=19ms° obtained for the distance and mass x] the upper limit shifts to d.κ724pc.," If this estimation is improved by $a^{\rm (dyn)}(R,\varphi=0.94)=-19\mbox{ms}^{-2}$ obtained for the distance and mass $[\ast]$, the upper limit shifts to $d < 724\mbox{pc}$."
+ From the functions M(7240.15.0.25-0.3.0.55) the upper mass limit is 4κ0.84+03M! ," From the functions ${\cal M}(724\:\mbox{pc},\varphi=0.15,0.25\mbox{-}0.3,0.55)$ the upper mass limit is ${\cal M} < 0.84\pm .03\:{\cal M}_\odot$ ."
+These upperlimits are in accordance with the values X., These upperlimits are in accordance with the values $[\ast]$.
+ However. they should. be treated. as indicative only. because they are sensitive to the actual value of OOOhoOLOnhyFOI.," However, they should be treated as indicative only, because they are sensitive to the actual value of ${\dot\vartheta},{\ddot\vartheta},\partial h_0/\partial t,
+\partial^2 h_0/\partial t^2$."
+ These values could be determined by numerical dillerentiations from P.ου in 0.91.xV5<0.96 when the OSA is not optimal to obtain all of them.," These values could be determined by numerical differentiations from $\vartheta,h_0$ in $0.91 \leq \varphi \leq 0.96$ when the QSAA is not optimal to obtain all of them."
+ Hence. he phase island ως=0.93-0.94 could not be included in the mass and distance determination because a very transient good ο.ΛΑ is embedded in a strongly shocked interval anc he mass .Vf appearing in g. plavs à secondary role in the dyvnanmies.," Hence, the phase island $\varphi=0.93\mbox{-}0.94$ could not be included in the mass and distance determination because a very transient good QSAA is embedded in a strongly shocked interval and the mass ${\cal M}$ appearing in $g_{\rm s}$ plays a secondary role in the dynamics."
+ For the sake of completeness. it must be mentioned that he pair (σι=0.15 and 4»=0.3) in (10)) gives d=1485 pe.," For the sake of completeness, it must be mentioned that the pair $(\varphi_1=0.15$ and $\varphi_2=0.3)$ in \ref{107a}) ) gives $d=1485$ pc."
+ Adding this value to those involved in. s]. d—(800-148) pe and M=(1.15205)ΝΤ. is obtained.," Adding this value to those involved in $[\ast]$, $d=(800\pm 148)$ pc and ${\cal M}=(1.15\pm .05){\cal M}_\odot$ is obtained."
+ This clistance is closer o d=900 pesuggested. by the parallax x=1.111.15 mas (Perryman1997)., This distance is closer to $d=900$ pcsuggested by the parallax $\pi=1.11\pm 1.15$ mas \citep{hipp1}.
+. However. this 1485 pe wasomitted because it is an outlier point above the 360 level. it exceeds the upper limit 724 pe and the inaccurate knowledge ofa! CRO15)aU?03} makes it uncertain.," However, this $1485$ pc wasomitted because it is an outlier point above the $3\sigma$ level, it exceeds the upper limit $724$ pc and the inaccurate knowledge of $a^{\rm (dyn)}(R,0.15)-a^{\rm (dyn)}(R,0.3)$ makes it uncertain."
+ Close to the shocked. phases (i.e. for y=01. 0.4. 0.45. 0.7-0.76. 0.84. 0.86. 0.91. 0.94. 0.97). negative stellar masses were obtained if a0(Roo)=0.," Close to the shocked phases (i.e. for $\varphi=0.1$, $0.4$, $0.45$, $0.7\mbox{-}0.76$, $0.84$, $0.86$, $0.91$, $0.94$, $0.97$ ), negative stellar masses were obtained if $a^{\rm (dyn)}(R,\varphi)=0$."
+" ""The. condition Midog)>0 gives an upper Limit for aCRus)<Oat these phases ancl shows clearly the limitations of treating an lh atmosphere in the QSAA."," The condition ${\cal M}(d,\varphi) > 0$ gives an upper limit for $a^{\rm (dyn)}(R,\varphi) < 0$ at these phases and shows clearly the limitations of treating an RR atmosphere in the QSAA."
+ In these phases. the characteristic value q|71., In these phases the characteristic value $\vert q\vert\approx 1$.
+" The sharp peaks q(/.,; and q(A.4=0.65)513.8 are produced by Qve£0tzxος η Le. the atmosphere is almost. completely. free of acceleration."," The sharp peaks $q(R,\varphi=0.5)=38.8$ and $q(R,\varphi=0.65)=-13.8$ are produced by $\partial v/\partial t\approx -v\partial v/\partial r$ , i.e. the atmosphere is almost completely free of acceleration."
+ q(I?=0.05)Lilt is a result from the poor representation of the density stratification by (7)). ic. the dvnamical term is large ing. οa?=(17.RT33.1)ms7 producing aU(ius)zmLaefal|eOviür).," $q(R,\varphi=0.05)=11.4$ is a result from the poor representation of the density stratification by \ref{2.110}) ), i.e. the dynamical term is large in $g_{\rm e}-g_{\rm s}-a^{\rm
+ (dyn)}=(44.7-8.7-33.1)\:\mbox{ms}^{-2}$ producing $a^{\rm (dyn)}(R,\varphi) \approx 11(\partial v/\partial t+v\partial
+v/\partial r)$."
+ This is the start of the descending branch in the light curve. when the atmosphere starts expanding rapidly.," This is the start of the descending branch in the light curve, when the atmosphere starts expanding rapidly."
+ The averaged. surface gravity for the whole pulsation evele is (logga)—-73. which is in excellent agreement with loggo=2.69.2.72 from ο and (BV(Rie: photometry (Siegel1982.. Liu&Janes 1990)). respectively.," The averaged surface gravity for the whole pulsation cycle is $\langle\log g_{\rm e}\rangle=2.73$, which is in excellent agreement with $\langle\log g_{\rm e}\rangle=2.69,2.72$ from $uvby$ and $UBV(RI)_C$ photometry \citealt{sieg1}, \citealt{liuj1}) ), respectively."
+ Phe average elfective. temperature of (4).=6778 Ix. ds significantly higher than (4.3=6433.6400.6490 Ix (Siegel. 1982.. Liu 1990. Bareza 2003)).," The averaged effective temperature of $\langle T_{\rm e}\rangle=6778$ K is significantly higher than $\langle T_{\rm e}\rangle=6433,6400,6490$ K \citealt{sieg1}, , \citealt{liuj1}, \citealt{barc1}) )."
+ The minimal.avcragec and maximal angular radii are (1.519.1.669.1.501) rad corresponding:to Ryin=(4.46.4.90.⋅− R.. respectively.," The minimal,averaged and maximal angular radii are $(1.519,1.669,1.801)\times 10^{-10}$ rad correspondingto $R_{\rm min}=(4.46,4.90,5.29)R_\odot$ , respectively."
+" “Phe following equilibrium. liniinosity5.29) ane ellective temperature (Carney.Strom&Jones1992). give the position of SU Dra in a theoretical Lertzsprung-Russe diagram (URD): Le,=A4xad(cy2)(45.9d 0.3)L.. T=[Laf/Agohd]y!(rqrabo2 20)Ilx."," The following equilibrium luminosity and effective temperature \citep{carn1} give the position of SU Dra in a theoretical Hertzsprung-Russell diagram (HRD): $L_{\rm eq}=4\pi\sigma
+d^2\langle\vartheta^2(\varphi) T_{\rm e}^4(\varphi)\rangle=(45.9\pm
+9.3)L_{\odot}$ , $T_{\rm eq}=\{L_{\rm eq}/4\pi\sigma
+[\langle\vartheta\rangle d]^2\}^{1/4} = \langle\vartheta^2T_{\rm
+ e}^4\rangle^{1/4} \langle\vartheta\rangle^{-1/2}=(6813\pm
+20)\mbox{K}$ ."
+ The magnitude averaged absolute brightness is Ady}=[0.68+23 mae.," The magnitude averaged absolute brightness is $\langle
+M_V \rangle=+0.68\pm .23$ mag."
+ From the point of view of radiative transfer. the approximation of the moclel atmospheres (Ixurucz could be applied. beyonddoubt. because. f+ holds for the whole pulsation evele of an It star.," From the point of view of radiative transfer, the plane-parallel approximation of the model atmospheres \citep{kuru1} could be applied beyonddoubt, because $h_0^{-1} \leq 0.03 R$ holds for the whole pulsation cycle of an RR star."
+ In favour of applying QSAA. semi-quantitative arguments," In favour of applying QSAA, semi-quantitative arguments"
+unclear region has mid-IR colours similar to those of late type galaxies 11.1. see Table 2)). although optically. it has very red colours sugecstive of an older stellar population (Struck et al. 1996)).,"nuclear region has mid-IR colours similar to those of late type galaxies 1.1, see Table \ref{phot}) ), although optically, it has very red colours suggestive of an older stellar population (Struck et al. \cite{struck}) )."
+ We have imientioned that Amram oet al. (1998]) , We have mentioned that Amram et al. \cite{amram}) )
+inve recentlv detected low-level Πα emission youn the Cartwheel ceuter. a result coufirming other iudepeudoeut observations (Iliedon priv.," have recently detected low-level $\alpha$ emission from the Cartwheel center, a result confirming other independent observations (Higdon priv."
+ conum)., comm.).
+ If one assunes that he Πα emissiou originates iu normal regions. could he observed LAW3 endssiou be explained in terms of wari dust heated by vouug stars?," If one assumes that the $\alpha$ emission originates in normal regions, could the observed LW3 emission be explained in terms of warm dust heated by young stars?"
+ We can estimate the amount of LW3 flux produced by star formation iu the micleus if we use the LW3 to Πα correlation found iu AI51 (Sauvage et al. 1996))., We can estimate the amount of LW3 flux produced by star formation in the nucleus if we use the LW3 to $\alpha$ correlation found in M51 (Sauvage et al. \cite{sauvage}) ).
+ Supposing that the Πα fiux detected by Αι et al. (1998..," Supposing that the $\alpha$ flux detected by Amram et al. \cite{amram},"
+ see Fig., see Fig.
+ 1). is distribute over the immer rug annuhts (area —170 aresec?) of the Cartwheel and that typically LW3/Tla — 30. we fix that which is of the same order as the observed LW3 enüssionu from the Cartwheel ceuter.," 1), is distributed over the inner ring annulus (area $\sim$ 170 $^{2}$ ) of the Cartwheel and that typically $\alpha$ $\sim$ 30, we find that, which is of the same order as the observed LW3 emission from the Cartwheel center."
+ Weuce. it appears that despite its ratler low-intensity. star formation activity from the Cartwhee center ds sufficieut o heat the dust and produce the observed wid-IR enüssion.," Hence, it appears that despite its rather low-intensity, star formation activity from the Cartwheel center is sufficient to heat the dust and produce the observed mid-IR emission."
+ Another possible source of dust heating iun the miclear region couk v0 attributed to the infall of gas clouds., Another possible source of dust heating in the nuclear region could be attributed to the infall of gas clouds.
+ This was proposed by Struck et al. (1996)), This was proposed by Struck et al. \cite{struck}) )
+ sed ou the morphooev of kiloparsec size. cometary-ike structures with blue huuinositfies iu the range 1.041079 +. detected in the edge of the iuncr rine.," based on the morphology of kiloparsec size, cometary-like structures with blue luminosities in the range $\times 10^{40}$ $^{-1}$, detected in the edge of the inner ring."
+ An order of magnitude calculation by the authors suggests hat the dissipation of the kinetic energy of the accreting clouds via shocks. would be —10/9 1. sufficient to eenerate a fraction of the observed blue luminosities.," An order of magnitude calculation by the authors suggests that the dissipation of the kinetic energy of the accreting clouds via shocks, would be $\sim$ $^{40}$ $^{-1}$, sufficient to generate a fraction of the observed blue luminosities."
+ We have obtained mid-IR ISOCAALD broad-band images of the Cartwheel eroup and comparing our data with our siuuple of normal aud active galaxies we were able to draw the following conclusious: 1) A huge scement of the outer ring is detected in the LW2 filter which is mainly dominated by thermally-spiked PATI cuission bands. while at longer wavelengths (LAV3 filter). where the cussion is primarily due to dust eraius in nearly thermal equilibiiun. the main source of enissiou originates from a single hot-spot in the rius associated with a particularly bright complex of reeious.," We have obtained mid-IR ISOCAM broad-band images of the Cartwheel group and comparing our data with our sample of normal and active galaxies we were able to draw the following conclusions: 1) A large segment of the outer ring is detected in the LW2 filter which is mainly dominated by thermally-spiked PAH emission bands, while at longer wavelengths (LW3 filter), where the emission is primarily due to dust grains in nearly thermal equilibrium, the main source of emission originates from a single hot-spot in the ring associated with a particularly bright complex of regions."
+ The hot-spot has an exceptional mid-IR broad-band diagnostic ratio LW3/LN2 of 5.2 which is among the lighes of any reeion in the CAAMACTTV sample aud is differcut TOU other regions of the rine., The hot-spot has an exceptional mid-IR broad-band diagnostic ratio LW3/LW2 of 5.2 which is among the highest of any region in the CAMACTIV sample and is different from other regions of the ring.
+ 2) A laree fraction of the mid-IR cussion is associated with the imuer ring and nucleus of the Cartwheel. in stark coutrast to that expected from optical cimission-line studies (where most of the lue enüssion originates from the outer ring).," 2) A large fraction of the mid-IR emission is associated with the inner ring and nucleus of the Cartwheel, in stark contrast to that expected from optical emission-line studies (where most of the line emission originates from the outer ring)."
+ Receutly faint Ta emission has been detected from the iuner ring aud it is possible that this nav be due to a low-level star formation which heats the erains., Recently faint $\alpha$ emission has been detected from the inner ring and it is possible that this may be due to a low-level star formation which heats the grains.
+ However. iu order to explain why the nuclear enissjon ds so powerful compared with the outer ring at ISOCAAT waveleneths. it secs that the dust of the outer ring must be spatially distributed very differently.," However, in order to explain why the nuclear emission is so powerful compared with the outer ring at ISOCAM wavelengths, it seems that the dust of the outer ring must be spatially distributed very differently."
+ One possibility is that the eraius in the outer ring experience a significantly diluted UW radiation field because they are lifte« out of the disk bv strouger stellar winds., One possibility is that the grains in the outer ring experience a significantly diluted UV radiation field because they are lifted out of the disk by stronger stellar winds.
+ Alternatively. the unclear regions iav be heated by a very cüffereu xocess than the outer. for example shock waves from iufaine clouds (Struck et al 1996).," Alternatively, the nuclear regions may be heated by a very different process than the outer, for example shock waves from infalling clouds (Struck et al \cite{struck}) )."
+ 3) The nüd-IBR cinission from the two conrpanious is typical for their Hubble type., 3) The mid-IR emission from the two companions is typical for their Hubble type.
+ Eveu rough our observations shed some more light to the properties of the hot dust in the Cartwheel galaxy. the a1nount and spatial distribution of the cold dus remain uncertain.," Even though our observations shed some more light to the properties of the hot dust in the Cartwheel galaxy, the amount and spatial distribution of the cold dust remain uncertain."
+ The 100445 TRAS flux of the galaxy is L6JJv but previous efforts to detect CO ποΓΙ (IIorellou 1995)) were unsuccessful. setting an upper lini to the Πω mass of 110° ," The $\mu m$ IRAS flux of the galaxy is Jy but previous efforts to detect CO emission (Horellou \cite{horellou}) ) were unsuccessful, setting an upper limit to the $_{2}$ mass of $^{9}$ $_{\odot}$."
+ Could this be explained by the low metallicity of the svstem. by its intrinsically low molecular gas coutent or simply bv the large distance of the Cartwheel?," Could this be explained by the low metallicity of the system, by its intrinsically low molecular gas content or simply by the large distance of the Cartwheel?"
+ Where is the peak of the spectral energy distribution iu this galaxy?, Where is the peak of the spectral energy distribution in this galaxy?
+ Deep sub-uu and nuu wave observations which ive scheduled iu the near future shoul enable us to address these questious., Deep sub-mm and mm wave observations which are scheduled in the near future should enable us to address these questions.
+noise.,noise.
+ ConsequentIy the velocity dispersion is lower (and in all our results subsonic) for the uniform shocks., Consequently the velocity dispersion is lower (and in all our results subsonic) for the uniform shocks.
+ Except. [or Fie., Except for Fig.
+ 6 though. we use the SPILL velocities. since these. are the velocities produced by the codo.," 6 though, we use the SPH velocities, since these are the velocities produced by the code."
+ We plot the velocity. dispersion size-scale relation for each of the simulations in Fig., We plot the velocity dispersion size-scale relation for each of the simulations in Fig.
+ 4. 5. 6. 7 and 8.," 4, 5, 6, 7 and 8."
+ Also shown is the 0x17 relation. coinciding with most observational results.," Also shown is the $\sigma \propto r^{0.5}$ relation, coinciding with most observational results."
+ In Fig., In Fig.
+ 4 and 5. we show the velocity size-scale relation for uniform: ancl clumpy initial eas clistributions. for the shock tube test ancl the sinusoidal potential.," 4 and 5, we show the velocity size-scale relation for uniform and clumpy initial gas distributions, for the shock tube test and the sinusoidal potential."
+ For the sinusoidal potential. the velocity dispersion. for the uniform shock dispersion. remains [lat anc subsonic (the sound. speed in all simulations is 0.3).," For the sinusoidal potential, the velocity dispersion for the uniform shock dispersion remains flat and subsonic (the sound speed in all simulations is 0.3)."
+ This is as expected. since for a uniform shock. the velocity of the shocked. gas should have zero velocity. dispersion.," This is as expected, since for a uniform shock, the velocity of the shocked gas should have zero velocity dispersion."
+ Again for the shock tube tests. the velocity. dispersions for the uniform shocks are relatively αι.," Again for the shock tube tests, the velocity dispersions for the uniform shocks are relatively flat."
+ By contrast the clumpsy shocks show an increasing velocity dispersion. with size-scale., By contrast the clumpy shocks show an increasing velocity dispersion with size-scale.
+ Phe velocity of gas in the shock depends on the amount of mass. it was encountered. (Section 3.3)., The velocity of gas in the shock depends on the amount of mass it has encountered (Section 3.3).
+ For the clumpy shock. gas entering the shock will encounter dilferent amounts of mass (c.g. where gas approaches another clump. or alternatively a relatively empty area) and a range of velocities are exhibited » the shocked. gas.," For the clumpy shock, gas entering the shock will encounter different amounts of mass (e.g. where gas approaches another clump, or alternatively a relatively empty area) and a range of velocities are exhibited by the shocked gas."
+ This range of velocities increases as he size-scale of the region increases., This range of velocities increases as the size-scale of the region increases.
+ At some size-scale. the region of gas will contain the full range of structure inherent in the initial distribution.," At some size-scale, the region of gas will contain the full range of structure inherent in the initial distribution."
+ The velocity size-scale relation hen remains relatively [lat for any further increase in size-scale., The velocity size-scale relation then remains relatively flat for any further increase in size-scale.
+ Fie., Fig.
+ 5 also compares the velocity. dispersion. relation when constant pressure boundaries are applied compared to using a dilluse hot medium., 5 also compares the velocity dispersion relation when constant pressure boundaries are applied compared to using a diffuse hot medium.
+ The average gradient [rom the 2 slopes is very. similar., The average gradient from the 2 slopes is very similar.
+ Although there is some cdillerence in the form of the velocity dispersion relation. the dvnamies appear to be dominated hy the cold. gas.," Although there is some difference in the form of the velocity dispersion relation, the dynamics appear to be dominated by the cold gas."
+ Fig., Fig.
+ 6 compares the velocity dispersion. caleulatecl using the actual and smoothed velocities. for a uniform and clumpy shock.," 6 compares the velocity dispersion calculated using the actual and smoothed velocities, for a uniform and clumpy shock."
+ Using the smoothecl velocities reduces the velocity. dispersion in the uniform shock by a factor of around. 2., Using the smoothed velocities reduces the velocity dispersion in the uniform shock by a factor of around 2.
+ The velocity, The velocity
+Infrared-dark clouds (RDCs) are a class of molecular elouds seen in extinction against the diffuse infrared emission of the Galactic plane.,Infrared-dark clouds (IRDCs) are a class of molecular clouds seen in extinction against the diffuse infrared emission of the Galactic plane.
+" About 2000 of these objects were discovered bv the Midcourse Space Experiment (MSX) satellite in spectral bands from 7—25, ", About 2000 of these objects were discovered by the Midcourse Space Experiment (MSX) satellite in spectral bands from $7-25\micron$ 
+"bv (e.g.. Frank 1995) Tir)e67001,/3.000.N)(/10R,)I?Ix. where 2, is the radius of the mass losing star. and. 7; is its surface temperature.","by (e.g., Frank 1995) $T(r) \simeq 670 (T_a/3,000 \K)(r/10 R_a)^{-1/2} \K$, where $R_a$ is the radius of the mass losing star, and $T_a$ is its surface temperature."
+" The corresponding sound speed is Cc301,/3.000IK).F?(r/10R,)kms.!."," The corresponding sound speed is $C \simeq 3 (T_a/3,000 \K)^{-1/2} (r/10 R_a)^{1/4} \km \s^{-1}$."
+ The wind Mach nuniber for most relevant cases is Machz5., The wind Mach number for most relevant cases is $\gtrsim 5$.
+ I find now the appropriate value of 5 for the problem at hand., I find now the appropriate value of $\gamma$ for the problem at hand.
+ The flow velocity from an AGDstaris e;cIOkms!., The flow velocity from an AGB star is $v_s \simeq 10 \km \s^{-1}$.
+ The flow is mostly neutral and dusty., The flow is mostly neutral and dusty.
+ The postshock temperature will be T;<(3/16)punye2/hEc200000./10kms.1)?IX. where pany is the mean mass per particle. and fis the Doltzmann constant.," The postshock temperature will be $T_s \leq (3/16)\mu m_H v_s^2/k \simeq 2000 (v_s/10 \km \s^{-1})^2 \K$, where $\mu m_H$ is the mean mass per particle, and $k$ is the Boltzmann constant."
+ For such a postshock temperature a large Iraction ol the dust and molecule are likely to survive the shock. and serve as cooling agents for the eas.," For such a postshock temperature a large fraction of the dust and molecule are likely to survive the shock, and serve as cooling agents for the gas."
+ Dhe preshock densitv is given bv The post shock density will be a few times higher. becoming even higher in the accretion column (Pogorelov et 22000).," The preshock density is given by The post shock density will be a few times higher, becoming even higher in the accretion column (Pogorelov et 2000)."
+ Using figure (11) of Woitke. INrügger. Sedlmavr (1996). I find the cooling time [rom T<8.000IX and density p>10!&em* to be τωE1νι.," Using figure (11) of Woitke, Krügger, Sedlmayr (1996), I find the cooling time from $T \lesssim 8,000 \K$ and density $\rho \gtrsim 10^{-16}\g \cm ^{-3}$ to be $\tau_{\rm cool} \lesssim 1 \yr$."
+ For a density of p10&em and a temperature of T<2.000IK. the cooling time is still short. Z4S8vr.," For a density of $\rho \gtrsim 10^{-19} \g \cm^{-3}$ and a temperature of $ T \lesssim 2,000 \K$, the cooling time is still short, $\tau_{\rm cool} \lesssim 3 \yr$."
+ The typical flow time of most of the accreted gas is ↕≺∢∪∐≺∢↥⋯⇂≼↲⊔⋯↥⊔∐↲↕↽≻∪⊳∖⇁↥⊳∖⊽∐∪≺∢↳↽↖≺↽↔↴≀↧⊔∖⇁↓⋟∪↕⋅⊔∐↲↕↽≻≀↧↴↕⋅≀↧↴↕∐≼↲∥↲↕⋅⊳∖⇁↕⋅≼↲↥≼↲∖↽≀↧↴∐↥↥∪≀↧↴≺∢≺∢↕⋅≼↲∐∪∐↓⋟↕⋅∪∐↓≀↧↴∖∖⊽↕∐≼⇂ ol upper-AGB stars at orbital separations where accretion disks are likely to be formed (see previous section) is radiative., The typical flow time of most of the accreted gas is I conclude that the postshock gas for the parameters relevant to accretion from a wind of upper-AGB stars at orbital separations where accretion disks are likely to be formed (see previous section) is radiative.
+ Namely. the gas is raclialively cooling «uite efficiently. such that the flow in the accretion column has a high Mach number. and the effective aciabatic index is ~1.," Namely, the gas is radiatively cooling quite efficiently, such that the flow in the accretion column has a high Mach number, and the effective adiabatic index is $\gamma \sim 1$."
+ The implication of the high Mach umber and an ellicient radiative cooling is that most ol the mass is being accreted in a dense and cold flow near the accretion linethe accretion column., The implication of the high Mach number and an efficient radiative cooling is that most of the mass is being accreted in a dense and cold flow near the accretion line–the accretion column.
+ The case for an isothermal flow with verv large Mach number. ie.. pressure is negligible. has a simple solution (e.g.. Lyttleton 1972).," The case for an isothermal flow with very large Mach number, i.e., pressure is negligible, has a simple solution (e.g., Lyttleton 1972)."
+ In (his flow all stream lines hit the accretion line., In this flow all stream lines hit the accretion line.
+ The total mass accretion rate per unit length on the accretion line is constant mM= TRacepoty. Where py and ry are the density and velocity of the unperturbed flow: here they will be taken as ry=ry. ancl py from equation (4).," The total mass accretion rate per unit length on the accretion line is constant $\dot m = \pi R_{\rm acc} \rho_0 v_0$ , where $\rho_0$ and $v_0$ are the density and velocity of the unperturbed flow; here they will be taken as $v_0=v_s$, and $\rho_0$ from equation (4)."
+ The mass and momentum conservation equations for the flow along the accretion line have infinite nmmber of solutions for the flow from the stagnation point outward. but only one solution from the stagnation point toward the accreting body (Lyttleton 19172: see his figure 1).," The mass and momentum conservation equations for the flow along the accretion line have infinite number of solutions for the flow from the stagnation point outward, but only one solution from the stagnation point toward the accreting body (Lyttleton 1972; see his figure 1)."
+ The equation lorthe, The equation forthe
+treat the evolution of a neutron star in quasi stationary approximation.,treat the evolution of a neutron star in quasi stationary approximation.
+ Though the partial transformation of rotational energv mto thermal energv may considerably change the cooling behaviour of a neutron star (see. e.9..VanRiper 1991)). and also the variation of space time ecolctry nieht have an effect on it. we stucly here. as a first step. he simplest case of coustaut angular velocity.," Though the partial transformation of rotational energy into thermal energy may considerably change the cooling behaviour of a neutron star (see, e.g.,\cite{VanRiper91b}) ), and also the variation of space time geometry might have an effect on it, we study here, as a first step, the simplest case of constant angular velocity."
+ The stationary. axisviuuetric. and asvinptotic fat metric iu quasi isotropic coordinates reads where the metric coefficients pnptt0) are functions of r aud 0 only.," The stationary, axisymmetric, and asymptotic flat metric in quasi isotropic coordinates reads where the metric coefficients $g_{\mu\nu}=g_{\mu\nu}(r,\theta)$ are functions of $r$ and $\theta$ only."
+ The metric coefficients are determined by the Einstein equation (ec=G 1) G=antl. and the energvanonmentuni conservation V-T=0.," The metric coefficients are determined by the Einstein equation $c=G=1$ ) ${\bf G} = 8\pi{\bf T}$, and the energy-momentum conservation ${\bf\nabla}\cdot{\bf T}=0$."
+ The obtained elliptic differcutial equations (Bonazzolactal. 1993)) are solved. via a finite differeuce scheme (Schaah 1998)) once. before the cooling simulation starts.," The obtained elliptic differential equations \cite{Bonazzola93a}) ) are solved via a finite difference scheme \cite{Schaab97a}) ) once, before the cooling simulation starts."
+ Iun the case of unifonu rotation. (0=coust.. considered here. the equations for thermal evolution are (Miyallesotal.1993.Schaal 1998)) where h denote the heat fiux 3-vector in the comoving frame. Cy the heat capacity. € the neutrino enmissivitv. aud & the heat conductivity.," In the case of uniform rotation, $\Omega={\rm const.}$, considered here, the equations for thermal evolution are \cite{Miralles93,Schaab97a}) ) where $\bf h$ denote the heat flux 3-vector in the comoving frame, $C_{\rm
+V}$ the heat capacity, $\epsilon$ the neutrino emissivity, and $\kappa$ the heat conductivity."
+ The partial radial and angular differentials are abbreviated by ὃ aud Oy. respectively.," The partial radial and angular differentials are abbreviated by $\partial_r$ and $\partial_\theta$, respectively."
+ Thermal equilibrium is described by T=coust..," Thermal equilibrium is described by $\tilde T={\rm
+const.}$."
+ At the surface of the neutron starthe heat flux fy aud h» is determined by the normal leat flux ὃν where R(0) is the r-coordinate of the surface., At the surface of the neutron starthe heat flux $h_1$ and $h_2$ is determined by the normal heat flux $h_{\rm N}$ where $R(\theta)$ is the $r$ -coordinate of the surface.
+ hx ds taken from a nonanagnetic photosphere model which describes the temperature gradient in the region between ο=Mec? and the stays surface (ee. Corcdauundsonetal.1983))., $h_{\rm N}$ is taken from a non-magnetic photosphere model which describes the temperature gradient in the region between $e=10^{10}\gccm$ and the star's surface (e.g. \cite{Gudmundson83}) ).
+ Iu. these models Px(0) depends on the teiiperature αἲ the density 10MecmP and ou the surface eravity The parabolic differential equations obtained after inserting Eqs.," In these models $h_{\rm
+N}(\theta)$ depends on the temperature at the density $e=10^{10}\gccm$ and on the surface gravity The parabolic differential equations obtained after inserting Eqs."
+ aud iuto Eq., and into Eq.
+ me solved via an implicit fuite difference scheme by using a alternating direction inuplicit method., are solved via an implicit finite difference scheme by using a alternating direction implicit method.
+ This yields a nou-luear equation svsteun which can be solved iteratively., This yields a non-linear equation system which can be solved iteratively.
+ The obtained lear equation svstenis have tridiagonal cocficicnt imatrices which can be inverted rather fast., The obtained linear equation systems have tridiagonal coefficient matrices which can be inverted rather fast.
+ The correctness of the two cdinensioual code was checked by comparing the outcome of it with simple. analytically solvable models and with the results of the oue dimensional code described by Schaah et al. (," The correctness of the two dimensional code was checked by comparing the outcome of it with simple, analytically solvable models and with the results of the one dimensional code described by Schaab et al. ("
+1996).,1996).
+ We consider a superfluid neutron star model basing ou the relativistic Tartree-Fock equation of state labelled RIIFS iu Uber et al. (, We consider a superfluid neutron star model basing on the relativistic Hartree-Fock equation of state labelled RHF8 in Huber et al. (
+1997). which accounts for hyperonic degrees of freedom.,"1997), which accounts for hyperonic degrees of freedom."
+ The slobal properties of uniformly rotating iunodoels with fixed eravitational mass Af=LOA. ancl aneular velocity Q- 0. 0.5. and O.99OK are summarized in Tab. L1...," The global properties of uniformly rotating models with fixed gravitational mass $M=1.5\,M_\odot$ and angular velocity $\Omega=0$ , $0.5$, and $0.99\,\Omega_{\rm K}$ are summarized in Tab. \ref{tab:models}."
+ On denotes the aN possible Isepler augular velocity. above which mass shedding sets in.," $\Omega_{\rm K}$ denotes the maximum possible Kepler angular velocity, above which mass shedding sets in."
+ All models allow for both the direct! uucleon Urea and for the direct. Lyperon Urea processes (cf. Prakashctal. 1992))., All models allow for both the direct nucleon Urca and for the direct hyperon Urca processes (cf. \cite{Prakash92}) ).
+ All direct Urea processes are suppressed by nucleou and lambda pairing below the respective critical cluperature (cf. Schaabetal. 1998))., All direct Urca processes are suppressed by nucleon and lambda pairing below the respective critical temperature (cf. \cite{Schaab98b}) ).
+ The ineredicuts to he cooling sinulations are similar to those discussed w Schaal et al. (, The ingredients to the cooling simulations are similar to those discussed by Schaab et al. (
+1996) in detail and are published ou he Web /imput.htial).,1996) in detail and are published on the Web /input.html).
+ Figeue 1 shows the evolution of the surface eniperature as iueasured by a distant observer., Figure \ref{fig:cool} shows the evolution of the surface temperature as measured by a distant observer.
+ It is oossible to distinguish three epochs of evolution., It is possible to distinguish three epochs of evolution.
+ In the first epoch f=LOO vr. large temperature gradieuts occur in the iuterior of the neutron star.," In the first epoch $t\lesssim 100$ yr, large temperature gradients occur in the interior of the neutron star."
+ Racial teniperature eracdients Were already found dn one dineusional sinulatious (sec. for example. Richardsonetal. 1982)).," Radial temperature gradients were already found in one dimensional simulations (see, for example, \cite{Richardson82}) )."
+ Tn rotating neutron stars. azimuthal temperature gradieuts. which cause transverse reat flow. fio4 0. exist. too.," In rotating neutron stars, azimuthal temperature gradients, which cause transverse heat flow, $h_2\neq 0$ , exist, too."
+ The polar temperature is by models rotating with 9=0.5Oi aud 0.99 Of. respectively (see also Fig. 2)).," The polar temperature is by models rotating with $\Omega=0.5\,\Omega_{\rm K}$ and $0.99\,\Omega_{\rm K}$ respectively (see also Fig. \ref{fig:deviation}) )."
+ After about 100 vr the cooling wave, After about 100 yr the cooling wave
+"Hence for Dmax the disk radius will be equal to Strictly speaking, for outbursts starting in the inner disk regions the (inside-out))radius in Eq. (3))","Hence for $L_{{\rm max}}$ the disk radius will be equal to Strictly speaking, for ) outbursts starting in the inner disk regions the radius in Eq. \ref{mdotcrit}) )"
+ corresponds to the distance reached by the outside-propagating heating front., corresponds to the distance reached by the outside-propagating heating front.
+ This distance can be shorter than the actual disk outer radius., This distance can be shorter than the actual disk outer radius.
+" In such case, the value given in Eq. "," In such case, the value given in Eq. \ref{rmax}) )"
+is only a lower limit for the disk radius., is only a lower limit for the disk radius.
+ Note also (5))that this assumes a non-irradiated disk., Note also that this assumes a non-irradiated disk.
+ Taking into account irradiation of the outer regions of the disk by X-rays from the inner regions always moves the critical radius further out1999)., Taking into account irradiation of the outer regions of the disk by X-rays from the inner regions always moves the critical radius further out.
+". Alternatively, an irradiation-dominated disk reprocessing passively Cz0.1 of the X-ray luminosity would still see its temperature fall below 8000 K (hydrogen ionization) only for radii greater than 2x10/? cm, similar to the above estimate."," Alternatively, an irradiation-dominated disk reprocessing passively ${\cal C} \approx 0.1$ of the X-ray luminosity would still see its temperature fall below 8000 K (hydrogen ionization) only for radii greater than $\times 10^{13}$ cm, similar to the above estimate."
+" So, again, Eq. (5))"," So, again, Eq. \ref{rmax}) )"
+ is only a lower limit on the disk radius., is only a lower limit on the disk radius.
+ From the Kepler's law a is with P4 being the orbital period (in days)., From the Kepler's law $a$ is with $P_{d}$ being the orbital period (in days).
+" Combining with Eq.(2)), the mean density of the companion p can be written in terms of the orbital period only A disk size Rp>5x1015 cm implies an orbital period = 23 days."," Combining with \ref{roche}) ), the mean density of the companion $\bar{\rho}$ can be written in terms of the orbital period only A disk size $R_{{\rm D}}\ga5\times10^{13}$ cm implies an orbital period $\ga$ 23 days."
+ The secondary star mean density for the hypothetical companion of HLX-1 gives 107gem? suggesting a red giant or a massive supergiant.," The secondary star mean density for the hypothetical companion of HLX-1 gives $10^{-4}\ {\rm g\, cm^{-3}}$ suggesting a red giant or a massive supergiant."
+" A priori, this is consistent with several possible evolutionary scenarios for HLX-13)."," A priori, this is consistent with several possible evolutionary scenarios for HLX-1."
+". The presumed outburst of HLX-1 does not look like a ""standard"" full-blown outburst of a black-hole transient low-mass X-ray binary.", The presumed outburst of HLX-1 does not look like a “standard” full-blown outburst of a black-hole transient low-mass X-ray binary.
+" In all such phenomena, after the outburst, the system declines (sometimes after one or two re-bounds) to a quiescent, very-low luminosity state."," In all such phenomena, after the outburst, the system declines (sometimes after one or two re-bounds) to a quiescent, very-low luminosity state."
+ According to the DIM the (non-constant) accretion rate is then everywhere lower than the critical value given by Eq. (3)) (, According to the DIM the (non-constant) accretion rate is then everywhere lower than the critical value given by Eq. \ref{mdotcrit}) ) (
+"in fact it is lower than the lower critical value Mz;,).",in fact it is lower than the lower critical value $\dot{M}_{{\rm crit}}^{-}$ ).
+ Assuming that the disk terminates at the ISCO Stable Circular Orbit) this implies a luminosity of (Innermost<2x10°° ss~!., Assuming that the disk terminates at the ISCO (Innermost Stable Circular Orbit) this implies a luminosity of $<2\times10^{30}$ $^{-1}$.
+" Even allowing for inner disk evaporation one can only increase this luminosity by 4 or 5, say, orders of magnitude2001),, still well below the observed ~2.6x104° ss-!."," Even allowing for inner disk evaporation one can only increase this luminosity by 4 or 5, say, orders of magnitude, still well below the observed $\sim2.6\times10^{40}$ $^{-1}$."
+ The moderate amplitude of the outburst in HLX-1 suggests that only part of the disk is involved that the DIM heating and cooling fronts propagate only in a restricted domain of the disk., The moderate amplitude of the outburst in HLX-1 suggests that only part of the disk is involved that the DIM heating and cooling fronts propagate only in a restricted domain of the disk.
+ Cooling fronts can be stalled when X-ray irradiation from the inner regions of the disk prevents cooling in the outer regions., Cooling fronts can be stalled when X-ray irradiation from the inner regions of the disk prevents cooling in the outer regions.
+ If the disk irradiation is directly tied to the mass accretion rate onto the compact object this does not prevent the disk from emptying but the decay light curve falls down exponentially2001).., If the disk irradiation is directly tied to the mass accretion rate onto the compact object this does not prevent the disk from emptying but the decay light curve falls down exponentially.
+ More complicated behavior can arise if X-ray irradiation is not directly tied to the inner mass accretion rate., More complicated behavior can arise if X-ray irradiation is not directly tied to the inner mass accretion rate.
+" Here, irradiation by a hot supergiant star located just outside of the disk can have an impact on its stability, independently of the mass accretion rate."," Here, irradiation by a hot supergiant star located just outside of the disk can have an impact on its stability, independently of the mass accretion rate."
+ Such a possibility has been advocated by in the case of outbursts of the (presumably), Such a possibility has been advocated by in the case of outbursts of the (presumably)
+km ! have higher metallicities («[Fe/H])=—1.29+0.07) than systems with Av100 km | «[Fe/H])=-1.53+ 0.18).,"km $^{-1}$ have higher metallicities $\langle [{\rm Fe/H]}\rangle = -1.29\pm 0.07$ ) than systems with $\Delta 
+{\rm v} < 100$ km $^{-1}$ $\langle [{\rm Fe/H]}\rangle = -1.53\pm 0.18$ )."
+ The KS test shows that the null-hypothesis that the metallicities of these two sub-samples of absorbers at z>2 are drawn from the same population can be rejected with confidence at c.]., The KS test shows that the null-hypothesis that the metallicities of these two sub-samples of absorbers at $z>2$ are drawn from the same population can be rejected with confidence at c.l.
+ This result is in line with the expectations of both the down-sizing picture and the velocity width as a proxy for mass also supported by the recent simulations of(2008)., This result is in line with the expectations of both the down-sizing picture and the velocity width as a proxy for mass also supported by the recent simulations of.
+. Neither sub-DLAs nor DLAs appear as the dominant featureof the sub-sample of absorbers at z>2 with Av>100 km ! or the sub-sample of absorbers at ς>2 with Av«100 km ': the former (the latter) sub-sample is composed of €19%)) sub-DLAs and (S1%)) DLAs., Neither sub-DLAs nor DLAs appear as the dominant featureof the sub-sample of absorbers at $z>2$ with $\Delta {\rm v} > 100$ km $^{-1}$ or the sub-sample of absorbers at $z>2$ with $\Delta {\rm v} < 100$ km $^{-1}$: the former (the latter) sub-sample is composed of ) sub-DLAs and ) DLAs.
+ This suggests that the M(H.) limit between and DLAs is not simply due to a distinction in mass., This suggests that the $N$ ) limit between sub-DLAs and DLAs is not simply due to a distinction in mass.
+ The objection to the massive galaxy hypothesis is the slope of the metallicity evolution of sub-DLAs which appears to be much steeper than that of luminous (massive) star-forming galaxies at |<z3DmainBodyCitationEnd2087]li08., The objection to the massive galaxy hypothesis is the slope of the metallicity evolution of sub-DLAs which appears to be much steeper than that of luminous (massive) star-forming galaxies at $1<z<3$.
+ Finally. the inconsistency in connecting sub-DLAs with massive galaxies is that they outnumber DLAs by a factor of up to 8. whereas massive galaxies are far less common than low mass galaxies.," Finally, the inconsistency in connecting sub-DLAs with massive galaxies is that they outnumber DLAs by a factor of up to 8, whereas massive galaxies are far less common than low mass galaxies."
+ Evenwhen assuming that only the most metal-rich DLAs arise in massive galaxies. at z«1.7 more than of sub-DLAs have [Zn/H] metallicities >2/5 solar and hence lead to a too large proportion of massive galaxies.," Evenwhen assuming that only the most metal-rich sub-DLAs arise in massive galaxies, at $z<1.7$ more than of sub-DLAs have [Zn/H] metallicities $> 2/5$ solar and hence lead to a too large proportion of massive galaxies."
+" Although massive gas-rich galaxies may have a higher absorption cross-section. have shown that in the local Universe the number density of DLAs is still dominated by sub-AZ, galaxies."," Although massive gas-rich galaxies may have a higher absorption cross-section, have shown that in the local Universe the number density of DLAs is still dominated by $M_{\star}$ galaxies."
+ So. how can the high metallicities of the low-redshift sub-DLAs be?," So, how can the high metallicities of the low-redshift sub-DLAs be?"
+? It has often been suggested that dust in the higher column density DLAs leads to an observed distribution of metallicities that is biased towards low values., It has often been suggested that dust in the higher column density DLAs leads to an observed distribution of metallicities that is biased towards low values.
+ Indeed. if dust extinction is a strong function of the metal column density as observed in interstellar medium clouds2004). the metal-rich DLAs may be more affected by dust selection than metal-rich sub-DLAs2008).," Indeed, if dust extinction is a strong function of the metal column density as observed in interstellar medium clouds, the metal-rich DLAs may be more affected by dust selection than metal-rich sub-DLAs."
+. Such an effect could make DLAs appear less metal-rich than sub-DLAs., Such an effect could make DLAs appear less metal-rich than sub-DLAs.
+ So far. however. there is no evidence at either high or low redshifts for a large number of high-extinction DLAs with elevated metallicities that would be required to explain the +0.7 dex difference in metallicity between the low-redshift sub-DLAs and DLAs.," So far, however, there is no evidence at either high or low redshifts for a large number of high-extinction DLAs with elevated metallicities that would be required to explain the $+0.7$ dex difference in metallicity between the low-redshift sub-DLAs and DLAs."
+ Most recently. found that dust depletions in 0.6«z<1.7 sub- and DLAs in the Complete Optical and Radio Absorption Line System (CORALS) survey are consistent with magnitude limited samples.," Most recently, found that dust depletions in $0.6<z<1.7$ sub-DLAs and DLAs in the Complete Optical and Radio Absorption Line System (CORALS) survey are consistent with magnitude limited samples."
+ Can large-scale physical differences account for the metallicity ditferences between low-redshift sub-DLAs and?, Can large-scale physical differences account for the metallicity differences between low-redshift sub-DLAs and?
+? Recently. have shown that for a given mass. galaxies are more metal-rich if they have smaller qalf-light radii or lower specific star formation rates.," Recently, have shown that for a given mass, galaxies are more metal-rich if they have smaller half-light radii or lower specific star formation rates."
+ have also shown that galaxies in dense environments have slightly higher metallicities., have also shown that galaxies in dense environments have slightly higher metallicities.
+ However. the metallicity otfsets for both effects are far too small (~0.05 dex) to explain the differences between sub-DLAs and DLAs we observe.," However, the metallicity offsets for both effects are far too small $\sim 0.05$ dex) to explain the differences between sub-DLAs and DLAs we observe."
+ What about photoionization??, What about photoionization?
+ The lower column densities of sub-DLAs may imply that the neutral hydrogen is yartially ionized. in which case the observed to ratio may not be a robust measure of the total [Fe/H] abundance.," The lower column densities of sub-DLAs may imply that the neutral hydrogen is partially ionized, in which case the observed to ratio may not be a robust measure of the total [Fe/H] abundance."
+ The magnitude of these corrections is quite sensitive to the ionizing background. leading to some variation amongst results in the literature2008).," The magnitude of these corrections is quite sensitive to the ionizing background, leading to some variation amongst results in the literature."
+ Substantial ionization corrections are. however. certainly needed in some DLAs.," Substantial ionization corrections are, however, certainly needed in some sub-DLAs."
+ The larger ditference between the median NCH 1) of sub-DLAs and DLAs at low than high and the change in the shape of the ionizing background between z~3 and z~|. also raises the possibility that the typical ionization corrections for sub-DLAs may be higher at z«1.7 than z>1.7.," The larger difference between the median $N$ ) of sub-DLAs and DLAs at low than high and the change in the shape of the ionizing background between $z\sim 3$ and $z\sim 1$, also raises the possibility that the typical ionization corrections for sub-DLAs may be higher at $z<1.7$ than $z>1.7$ ."
+ However. there is currently no strong evidence to support this as an explanation for the much steeper evolution of sub-DLA metallicities.," However, there is currently no strong evidence to support this as an explanation for the much steeper evolution of sub-DLA metallicities."
+ In general. ionization corrections in sub-DLAs are similar when calculated with a Haardt-Madau spectrum at both z~3 and z~OS (Milutinovie. private communication) and found corrections that are typically « 0.2 dex for their low-redshift sub-DLAs.," In general, ionization corrections in sub-DLAs are similar when calculated with a Haardt-Madau spectrum at both $z\sim 3$ and $z\sim 0.5$ (Milutinovic, private communication) and found corrections that are typically $<$ 0.2 dex for their low-redshift sub-DLAs."
+ Having ruled out physically related sources of bias (dust. environment. and ionization corrections). we consider svstematics associated with the selection and analysis of low-redshift DLAs.," Having ruled out physically related sources of bias (dust, environment, and ionization corrections), we consider systematics associated with the selection and analysis of low-redshift sub-DLAs."
+ This is additionally motivated by our finding that the sub-DLAs only show significant enhancement in [Fe/H] at redshifts that require HST observations of the Lya line., This is additionally motivated by our finding that the sub-DLAs only show significant enhancement in [Fe/H] at redshifts that require HST observations of the $\alpha$ line.
+ have investigated possible systematic uncertainties in the column density determinations of sub-DLAs from low-resolution spectra in comparison to resolution spectra., have investigated possible systematic uncertainties in the column density determinations of sub-DLAs from low-resolution spectra in comparison to high-resolution spectra.
+ They have found that NCH 1) derived from low-resolution spectra tend to be over-estimated by typically 0.1 dex (up to 0.3 dex). but in the wrong direction to explain the high metallicities observed in the low-redshift sub-DLAs.," They have found that $N$ ) derived from low-resolution spectra tend to be over-estimated by typically 0.1 dex (up to 0.3 dex), but in the wrong direction to explain the high metallicities observed in the low-redshift sub-DLAs."
+ Next. we consider the process through which the majority of our low-redshift sub-DLAs were identified: the equivalent width (EW) selection.," Next, we consider the process through which the majority of our low-redshift sub-DLAs were identified: the equivalent width (EW) selection."
+ Most of the low-redshift sub-DLAs in both this study and the literature come from compilations of sub-DLAs that have been obtained as a “by-product” of low-redshift DLA searches., Most of the low-redshift sub-DLAs in both this study and the literature come from compilations of sub-DLAs that have been obtained as a `by-product' of low-redshift DLA searches.
+ These DLA surveys have used metal lines to pre-select candidate DLAs from ground-based spectra for HST follow-up., These DLA surveys have used metal lines to pre-select candidate DLAs from ground-based spectra for HST follow-up.
+ have found that there are no DLAs in their unbiased sample (.e. excluding previously known 21 em absorbers} with rest-frame Mgu.t2796 EWs lower than 0.6 A., have found that there are no DLAs in their unbiased sample (i.e. excluding previously known 21 cm absorbers) with rest-frame $\lambda$ 2796 EWs lower than 0.6 .
+. Pre-selecting candidate DLAs based on their EWs therefore appears to be a reasonable selection strategy of DLAs2004)., Pre-selecting candidate DLAs based on their EWs therefore appears to be a reasonable selection strategy of DLAs.
+. However. at lower column densities. an increasing fraction of QAL systems have 12796 EWs «0.6citeprao06..," However, at lower column densities, an increasing fraction of QAL systems have $\lambda$ 2796 EWs $< 0.6$."
+ As a consequence. a high EW cut does not select all sub-DLAs in an unbiased way.," As a consequence, a high EW cut does not select all sub-DLAs in an unbiased way."
+ In Fig., In Fig.
+ 3. we show the rest-frame Mgi12796 EWs and metallicities for sub-DLAs and DLAs at z«1.7 of our sample and the sample of (2009)., \ref{EW-metallicity} we show the rest-frame $\lambda$ 2796 EWs and metallicities for sub-DLAs and DLAs at $z<1.7$ of our sample and the sample of .
+. Clearly. sub-DLAs are metal-rich for their EWs compared with the DLAs. and this despite that the 1.12796 EW distributions of sub-DLAs and DLAsat z«1.7 are statistically consistent (see their cumulative functions: KS c.l.," Clearly, sub-DLAs are metal-rich for their EWs compared with the DLAs, and this despite that the $\lambda$ 2796 EW distributions of sub-DLAs and DLAsat $z<1.7$ are statistically consistent (see their cumulative functions; KS c.l."
+ to reject the null-hypothesis = 105€)., to reject the null-hypothesis $= 10$ ).
+ This offset can be explained in terms of Kinematics., This offset can be explained in terms of kinematics.
+ and showed that dow-redshift) sub-DLAs have higher, and showed that (low-redshift) sub-DLAs have higher
+"Having obtained the luminosity function and characterized some of the physical properties of the current high-z candidate galaxies, we now turn to an analysis of their role in the reionization process.","Having obtained the luminosity function and characterized some of the physical properties of the current $z$ candidate galaxies, we now turn to an analysis of their role in the reionization process."
+" The rate of ionizing photons from the j-th galaxy is given by where Φ (QU) is the ionizing photon flux for Pop II (Pop III) stars, which of course is a function of the IMF, metallicity and stellar age of any given galaxy."," The rate of ionizing photons from the $j$ -th galaxy is given by where $Q^{\rm II}$ $Q^{\rm III}$ ) is the ionizing photon flux for Pop II (Pop III) stars, which of course is a function of the IMF, metallicity and stellar age of any given galaxy."
+ Finally [εις is the escape fraction of ionizing photons., Finally $f_{esc}$ is the escape fraction of ionizing photons.
+" For simplicity, we assume a redshift-independent value of f... for the two stellar populations."," For simplicity, we assume a redshift-independent value of $f_{esc}$ for the two stellar populations."
+" The ionization rate provided by galaxies in the simulation box per unit comoving volume at redshift z, Nion(z), is then given by the sum over all galaxies at that redshift divided by the volume of the simulation."," The ionization rate provided by galaxies in the simulation box per unit comoving volume at redshift $z$, $\dot{N}_{ion}(z)$, is then given by the sum over all galaxies at that redshift divided by the volume of the simulation."
+ 'The actual ionization rate must also include galaxies that are too rare to be caught in our relatively small simulation volume., The actual ionization rate must also include galaxies that are too rare to be caught in our relatively small simulation volume.
+" To account for this correction, we add the ionization rate due to bright galaxies by integrating the observed LF (or the upper limits) from the luminosity of the brightest galaxy in the output down to very low magnitudes (Muv= —25)."," To account for this correction, we add the ionization rate due to bright galaxies by integrating the observed LF (or the upper limits) from the luminosity of the brightest galaxy in the output down to very low magnitudes $M_{UV}=-25$ )."
+ The ionizing photon flux is obtained using a SED derived assuming that rare galaxies have the same stellar age and metallicity as the brightest simulated one., The ionizing photon flux is obtained using a SED derived assuming that rare galaxies have the same stellar age and metallicity as the brightest simulated one.
+ The extra contribution of the unaccounted bright-end of the LF is found to be at most 1096 of the total ionizing photon emission., The extra contribution of the unaccounted bright-end of the LF is found to be at most $10$ of the total ionizing photon emission.
+" The evolution of the total specific ionizationrate, Nion, is plotted in the top panel of Fig."," The evolution of the total specific ionizationrate, $\dot{N}_{ion}$, is plotted in the top panel of Fig."
+" 8 (f...=0.2) along with the same rate due to galaxies detected in the HUDF, NE. or detectable by JWST, N7,,."," \ref{fig:reion} $f_{esc}=0.2$ ) along with the same rate due to galaxies detected in the HUDF, $\dot{N}_{ion}^H$, or detectable by JWST, $\dot{N}_{ion}^J$."
+" The ratios between (NE, Ni.) and Nion are also shown for clarity in the bottom panel of the same Figure."," The ratios between $\dot{N}_{ion}^H$, $\dot{N}_{ion}^J$ ) and $\dot{N}_{ion}$ are also shown for clarity in the bottom panel of the same Figure."
+ HST is now resolving the sources that provide z1/3 of the ionizing photon budget at z=5 and ~20% at z=7—7.5., HST is now resolving the sources that provide $\approx 1/3$ of the ionizing photon budget at $z=5$ and $\sim 20$ at $z=7-7.5$.
+ This results is consistent with the limits set by the observed very steep faint-end slope of thez —7—8 LF (Bouwens et al., This results is consistent with the limits set by the observed very steep faint-end slope of the $z=7-8$ LF (Bouwens et al.
+ 2010b) and with the estimate of recent semi-analytical models (Choudhury et al., 2010b) and with the estimate of recent semi-analytical models (Choudhury et al.
+" 2008, Trenti et al."," 2008, Trenti et al."
+ 2010)., 2010).
+" At the sensitivity limit of JWST, it will be possible to detect the bulk of ionizing sources up to z~7.3, but at higher redshifts most of the ionizing photons will still be produced by sources that are too faint to be detected even by JWST."," At the sensitivity limit of JWST, it will be possible to detect the bulk of ionizing sources up to $z\sim 7.3$, but at higher redshifts most of the ionizing photons will still be produced by sources that are too faint to be detected even by JWST."
+" The total ionization rate density Nion(z) should then be compared with therecombination rate density of the IGM, N.««(z), given by (e.g. Madau, Haardt Rees 1999) where (nz) is the mean comoving hydrogen density in the Universe and (trec) is the volume-averaged recombination time for ionized hydrogen with an effective Hy clumping factor Cur;=(niugj)/(nnmi)."," The total ionization rate density $\dot{N}_{ion}(z)$ should then be compared with therecombination rate density of the IGM, $\dot{N}_{rec}(z)$, given by (e.g. Madau, Haardt Rees 1999) where $\langle n_H \rangle$ is the mean comoving hydrogen density in the Universe and $\langle t_{rec} \rangle$ is the volume-averaged recombination time for ionized hydrogen with an effective $_{\rm II}$ clumping factor $C_{HII}=\langle n^2_{HII}\rangle/\langle n_{HII}\rangle^2$."
+ The recombination rate density is shown —in the top panel of Fig., The recombination rate density is shown in the top panel of Fig.
+ 8 with dotted lines for different value of the clumping factor Curr., \ref{fig:reion} with dotted lines for different value of the clumping factor $C_{HII}$.
+" For fese=0.2 and Cyr;=10, the balance between ionization and recombination is obtained at z~6.8."," For $f_{esc}=0.2$ and $C_{HII}=10$, the balance between ionization and recombination is obtained at $z\sim 6.8$."
+ For fesc=0.1 Nion=View at z=6 assuming Cyr10., For $f_{esc}=0.1$ $\dot{N}_{ion}=\dot{N}_{rec}$ at $z=6$ assuming $C_{HII}=10$.
+" Another piece of useful information than can be extracted from the simulation outputs is the relative fraction of normal (Pop II) and massive, metal-free (Pop III) stars."," Another piece of useful information than can be extracted from the simulation outputs is the relative fraction of normal (Pop II) and massive, metal-free (Pop III) stars."
+ There are several questions to which we can provide quantitative answer: (i) do some of the current candidates contain Pop III stars? (, There are several questions to which we can provide quantitative answer: (i) do some of the current candidates contain Pop III stars? (
+"ii) in that case, what fraction of their UV luminosity is powered by them? (","ii) in that case, what fraction of their UV luminosity is powered by them? ("
+iii) how is this fraction dependent on their Muv luminosity? (,iii) how is this fraction dependent on their $M_{UV}$ luminosity? (
+iv) is there a clear observational signature imprinted by Pop III stars?,iv) is there a clear observational signature imprinted by Pop III stars?
+" The answer to the first question is straightforward: having analyzed the stellar populations of the simulated galaxies present at four observationally relevant redshifts, z=(5,6,8,10), we find that a fraction 0.07-0.19 (depending on z) of the galaxies contain at least some Pop III stars."," The answer to the first question is straightforward: having analyzed the stellar populations of the simulated galaxies present at four observationally relevant redshifts, $z=(5, 6, 8, 10)$, we find that a fraction 0.07-0.19 (depending on $z$ ) of the galaxies contain at least some Pop III stars."
+ We should not emphasize too much on the exact values of this Pop III/Pop II galaxy ratio as fluctuations in galaxy mass and star formation rate might introduce a very large dispersion., We should not emphasize too much on the exact values of this Pop III/Pop II galaxy ratio as fluctuations in galaxy mass and star formation rate might introduce a very large dispersion.
+ The most robust physical quantity to understand the relative importance of the two populations is the ratio of Pop III-to-Pop II star formation rates which is a decreasing function of time (see Fig., The most robust physical quantity to understand the relative importance of the two populations is the ratio of Pop III-to-Pop II star formation rates which is a decreasing function of time (see Fig.
+ 1 of TFS07) never exceeding 10? below z— 10., 1 of TFS07) never exceeding $10^{-3}$ below $z=10$ .
+" More relevant is question (ii above, whose answer can be obtained by inspecting Fig. 9,,"," More relevant is question (ii) above, whose answer can be obtained by inspecting Fig. \ref{fig:Pop III}, ,"
+" showing the ratio,"," showing the ratio,"
+ To measure the radial velocities of the four double clegenerates: WD1022|050.373... WDIS24|040 and W1D)2082|ISS. we used least squares fitting of a mocel line profile.," To measure the radial velocities of the four double degenerates: WD1022+050, WD1824+040 and WD2032+188, we used least squares fitting of a model line profile."
+ Ehe model line profile is the summation of three Gaussian profiles with different widths and depths., The model line profile is the summation of three Gaussian profiles with different widths and depths.
+ For any eiven star. the widths. and epths of the Caussians are optimised and then held. fixed. while their velocity. ollsets from the rest wavelengths of the lines in question are fitted separately for each spectrum: see Maxted. Marsh Moran (2000c) for further details of this procedure.," For any given star, the widths and depths of the Gaussians are optimised and then held fixed while their velocity offsets from the rest wavelengths of the lines in question are fitted separately for each spectrum; see Maxted, Marsh Moran \shortcite{m00c} for further details of this procedure."
+" Once the radial velocities for cach system were known (see Table 2)) we used a coating mean"" periodogram to determine the periods of our targets (e.g. Cumming. Marcy Butler 1999)."," Once the radial velocities for each system were known (see Table \ref{results:rv:tab2}) ) we used a “floating mean” periodogram to determine the periods of our targets (e.g. Cumming, Marcy Butler 1999)."
+ Phe method consists in fitting the data with a model composed of a sinusoid plus a constant of the form: where f is the frequency and {15 the observation time., The method consists in fitting the data with a model composed of a sinusoid plus a constant of the form: where $f$ is the frequency and $t$ is the observation time.
+ The kev point is that the systemic velocity is fitted at the same time as A and fy., The key point is that the systemic velocity is fitted at the same time as $K$ and $t_0$.
+ This corrects a failing of the well-known Lomb-Searele (Lomb1976:Scarele1982) periocogram which starts by subtracting the mean of the data and then fits a plain sinusoid: this is not the best approach for small numbers of points.," This corrects a failing of the well-known Lomb-Scargle \cite{l76,s82} periodogram which starts by subtracting the mean of the data and then fits a plain sinusoid; this is not the best approach for small numbers of points."
+ We obtained the 4? of the fit as a Function of f and then identified. minima in this function., We obtained the $\chi^2$ of the fit as a function of $f$ and then identified minima in this function.
+ ‘Table 2. gives a list of the orbital parameters derived for each DD binary star., Table \ref{res:rv:tab1} gives a list of the orbital parameters derived for each DD binary star.
+ Phe orbital period of the second best alias. is. also given.. along with. the differencemen in. V72 between the two best. periods fourο," The orbital period of the second best alias is also given, along with the difference in $\chi^2$ between the two best periods found."
+ The large dilference in X7 indicates that the second best aliases are not plausible., The large difference in $\chi^2$ indicates that the second best aliases are not plausible.
+ Phe resulting radial velocity. curves (folded. in the orbital period). are presented in Fig., The resulting radial velocity curves (folded in the orbital period) are presented in Fig.
+ 2. and the corresponding perioclograms (A7 versus orbital frequency) in Fig. 3.., \ref{res:rv:rvWD} and the corresponding periodograms $\chi^2$ versus orbital frequency) in Fig. \ref{res:rv:pgramWD}.
+ Each panel in the periodogram includes a blow up of the region in frecqucney where the minimum X7 is found., Each panel in the periodogram includes a blow up of the region in frequency where the minimum $\chi^2$ is found.
+ In each case. we compute the level of systematic uncertainty that when added: in quadrature to our raw error estimates gives a reduced 47~1.," In each case, we compute the level of systematic uncertainty that when added in quadrature to our raw error estimates gives a reduced $\chi^2
+\sim 1$."
+ By doing this we are considering the un-accounteck sources of error such as true variability of the star or slit-lilline errors that cause the poor fits of a few stars., By doing this we are considering the un-accounted sources of error such as true variability of the star or slit-filling errors that cause the poor fits of a few stars.
+ Such errors are unlikely to be correlated. with either the orbit or with the statistical errors we estimate. and therefore we add a fixed quantity in cuacirature with our statistical errors as opposed to applying a simple multiplicative scaling to them.," Such errors are unlikely to be correlated with either the orbit or with the statistical errors we estimate, and therefore we add a fixed quantity in quadrature with our statistical errors as opposed to applying a simple multiplicative scaling to them."
+" In all cases we use a minimum value of 2kms+ corresponding to 1/10"" of a"," In all cases we use a minimum value of $2\,{\rm
+km}\,{\rm s}^{-1}$ corresponding to $^{\rm th}$ of a"
+discrepancy is probably due to the uuusual streneth of CTI] in this source. rather than HeII being too weak. as it wmielt sccm from the diagnostic diagrams We have also studied the NVAL2lO cussion.,"discrepancy is probably due to the unusual strength of CIII] in this source, rather than HeII being too weak, as it might seem from the diagnostic diagrams 	We have also studied the $\lambda$ 1240 emission."
+ This line was not often detected iu earlier spectra of WzRC because of limited S/N {see van Ojik 1995))., This line was not often detected in earlier spectra of HzRG because of limited S/N (see van Ojik \cite{ojik95}) ).
+ Iun our sanall sample. ouly MRC2025-218 and SAINT J02399-0136 have detectable NV. cunission.," In our small sample, only MRC2025-218 and SMM J02399-0136 have detectable NV emission."
+ We present in Fig., We present in Fig.
+ & the diagram ΠΟΠ NV/CTV., \ref{Fig8} the diagram NV/HeII NV/CIV.
+ Quasars define a tieht correlation im this diagram (Ibuuaun Ferland 1993)) which is represented as a inclined line., Quasars define a tight correlation in this diagram (Hamann Ferland \cite{hamm93}) ) which is represented as a inclined line.
+ Fosbury ot al. (1998..1999))," Fosbury et al. \cite{fos98}, \cite{fos99}) )"
+ showed that Ην follow a parallel correlation to the one defined by quasars., showed that HzRG follow a parallel correlation to the one defined by quasars.
+ This is also shown in the diagram., This is also shown in the diagram.
+ We have plotted the position of SAIALIO2399-0136. LI and MIRC2025-218.," We have plotted the position of SMMJ02399-0136, L1 and MRC2025-218."
+ Tuterestinely. the NV/CTV. and NV/Uell line ratios measured iu these sources are the largest observed for ITZRC:.," Interestingly, the NV/CIV and NV/HeII line ratios measured in these sources are the largest observed for HzRG."
+ AIRC2025-218 lies at the top of the correlation defined by ITZRC:. while L1 lies ou the quasars correlation and also occupies the position of the largest values for the NV line ratios.," MRC2025-218 lies at the top of the rrelation defined by HzRG, while L1 lies on the quasars correlation and also occupies the position of the largest values for the NV line ratios."
+ Both standavd ACN photoionization models (with solar abundances and deusitv o «100 ) and shock models are unable to reproduce the position of the objects iu this diagram., Both standard AGN photoionization models (with solar abundances and density $n\leq$ 100 $^{-3}$ ) and shock models are unable to reproduce the position of the objects in this diagram.
+"uncorrelated.g(x)=05,07(x). where σι is theres noise.","uncorrelated,$\eta({\mathbf
+x}) = \sigma_{n}^{2} \delta^{(2)}({\mathbf x})$, where $\sigma_{n}$ is the noise."
+ As a result. the covariance matrix reduces to where we have used the orthonormalitv of the basis functions (Eq. 4]].," As a result, the covariance matrix reduces to where we have used the orthonormality of the basis functions (Eq. \ref{eq:orthonorm}] ])."
+ 1n this case. the covariance matrix is thus iagonal. so that each cocllicicnt is statistically independent.," In this case, the covariance matrix is thus diagonal, so that each coefficient is statistically independent."
+ Moreover. the diagonal elements are all equal.," Moreover, the diagonal elements are all equal."
+ Uneorrelated noise thus populates cach coefficient: equallv. and is thus “white” as in the case of Fourier transforms.," Uncorrelated noise thus populates each coefficient equally, and is thus “white” as in the case of Fourier transforms."
+ In the case of spatially correlated but. homogeneous noise. the noise correlation function is only a function of separation and can thus be written as g(xx’).," In the case of spatially correlated but homogeneous noise, the noise correlation function is only a function of separation and can thus be written as $\eta({\mathbf x} - {\mathbf x}')$."
+ As a result. Equation (39)) reduces to the integral of a convolution and can thus be written svmbolically as in the notation of Equation (45)).," As a result, Equation \ref{eq:cov_general}) ) reduces to the integral of a convolution and can thus be written symbolically as in the notation of Equation \ref{eq:cnml_notation}) )."
+ X convenient way to evaluate this is to decompose (x) itself into basis functions and then to use the results of €refeonvolution below.," A convenient way to evaluate this is to decompose $\eta({\mathbf x})$ itself into basis functions and then to use the results of \\ref{convolution}
+ below."
+ Spatial correlations in the noise thus produces correlations in the coellicients., Spatial correlations in the noise thus produces correlations in the coefficients.
+ Another case of practical interest is that in which the noise is dominated by Poisson shot noise., Another case of practical interest is that in which the noise is dominated by Poisson shot noise.
+" I£ the intensities are measured in units of photon counts. the noise correlation function is then (x.κ)=fix)κ.x"")."," If the intensities are measured in units of photon counts, the noise correlation function is then $\eta({\mathbf x},{\mathbf x}') = f({\mathbf x}) \delta^{(2)}({\mathbf
+x}-{\mathbf x}')$."
+ Ns a result. the COVALLANCE nialrix 15 ii . ∖∖⋎↓↥∢⊾↓⋅⋖⋅∠↗⋟⋣≺⊽⋟↕⊽⋯≼⊰⋡⊰⋡∩↓≱∖↥↓↕⋖⋅⇀↗≻−↓≻↓⋅⋯⊔∐∼↿↓⊔∣∢⊾⋏∙≟↓⋅⋜↧↓∠⇂∢⊾∐⋯⊾∠⇂. . ⋅ ⊲↓⊔∣⊲⇀⊲⊏↥⊔⋜⊔⊲↓∪↓⊔∖≟⊤∩∣⋡∢⋅↓∪∖∖⋎⊳⋜⋯∠⇂∖∖⋎↓↥↕≼∼↓↕↕≻⋖⋅∖⇁⋜↧↓⋯⋯⋅∠⇂⋜⋯⋜↧↓∙∖⇁↿⊲⊔∼⋜↧∐∙∖⇁ ⊲↓⊔↓↴⋜↧↓≻⋖⋅↓⋅∐⊳↓⊔↿↓↕⊲↓⊳∖≼∼⋜," As a result, the covariance matrix is where $B^{(3)}_{{\mathbf k},{\mathbf l},{\mathbf
+m}}(\beta,\beta,\beta)$ is the 3-product integral defined in Equation \ref{eq:b3}) ) below, and which is evaluated analytically in Paper II."
+↧⊳∖⋖⋅∥⋏∙≟∥⊀↓⊔⊳⇂↓↥⋖⊾≼⇍∪∖⇁⋜⊔⋰↓⋜⋯≼⇍⋖⊾≼⇍⋯⋅∐↕≼⋰⊓⋅⊔↥⋠↓≱∖ made non-diagonal by the noise correlation. but is casily caleulable analytically.," In this case again, the covariance coefficient is made non-diagonal by the noise correlation, but is easily calculable analytically."
+ We now show how shapelets behave under. convolutions. an operation which often occurs in practice (ce.," We now show how shapelets behave under convolutions, an operation which often occurs in practice (eg."
+ under the action of PSE. seeing. smoothing. ete).," under the action of PSF, seeing, smoothing, etc)."
+ We start. by considering convolution by a general kernel in l-cimensions. and then study the special case of smoothing by a gaussian.," We start by considering convolution by a general kernel in 1-dimensions, and then study the special case of smoothing by a gaussian."
+ Finally. we treat the 2-cimensional case. ancl illustrate the results with the example of an LIST galaxy image.," Finally, we treat the 2-dimensional case, and illustrate the results with the example of an HST galaxy image."
+ Let us first. consider the convolution of two arbitrary 1-dimensional functions fe) and gGe)., Let us first consider the convolution of two arbitrary 1-dimensional functions $f(x)$ and $g(x)$.
+ Their convolution hr) can be written as Each function can be decomposecl into our basis functions with scales o. 3 and 5.," Their convolution $h(x)$ can be written as Each function can be decomposed into our basis functions with scales $\alpha$ , $\beta$ and $\gamma$."
+ ποσο scales are chosen to be most convenient in each case., These scales are chosen to be most convenient in each case.
+ The cocllicients are then fu=Πταὃν gy={μιεν hy—Gass|f).," The coefficients are then $f_{n} \equiv \langle
+n;\alpha | f \rangle$, $g_{n} \equiv \langle n;\beta | g \rangle$, $h_{n} \equiv \langle n;\gamma | h \rangle$."
+" Our aim is to find an expression which relates /,, to fi, and qg,.", Our aim is to find an expression which relates $h_{n}$ to $f_{n}$ and $g_{n}$.
+ Since convolution is a bi-linear operation. this relation will be of the form where the convolution tensor can be written svmbolicallv as and is a function. of the scale. lengths.," Since convolution is a bi-linear operation, this relation will be of the form where the convolution tensor can be written symbolically as and is a function of the scale lengths."
+ Using the properties of the basis functions under Fourier transforms (Eq. 9].," Using the properties of the basis functions under Fourier transforms (Eq. \ref{eq:B_tilde}] ]),"
+ it is easy to show that the convolution tensor is eiven by where the 3-product integral is Be!iudανα.(αι). ds defined. as As we show in Paper LH. this integral can be easily evaluated analvtically with a recurrence relation.," it is easy to show that the convolution tensor is given by where the 3-product integral is $B^{(3)}_{nml}(a_1,a_2,a_3)$ is defined as As we show in Paper II, this integral can be easily evaluated analytically with a recurrence relation."
+ The special case consisting of smoothing by a gaussian is useful in practice., The special case consisting of smoothing by a gaussian is useful in practice.
+ In this case. we Let which is normalised so that {dvgir)= 1.," In this case, we let which is normalised so that $\int dx~g(x)=1$ ."
+ We ean then write the coellicients for the smoothed function hr) as where μία.d)=Σο”...Nar ds the smoothing matrix.," We can then write the coefficients for the smoothed function $h(x)$ as where $G_{nm}(\gamma,\alpha,\beta) = \sum_{l}
+C_{nml}(\gamma,\alpha,\beta) g_{l}$ is the smoothing matrix."
+" Phe gaussian gor) can be thought as a (non-normalised) m=O shapelet state of amplitude gy=(0:ή) so that Gai,=Conogo."," The gaussian $g(x)$ can be thought as a (non-normalised) $n=0$ shapelet state of amplitude $g_{0}=\langle 0;\beta | g \rangle$ , so that $G_{nm} = C_{nm0} g_{0}$."
+" Using the generating function for Lermite polynomials. one can show that. for the natural choice of 57=a72|377. the smoothing matrix 1s given by formon>0 and even (C,,,,, vanishes otherwise). and where awσα7|37, "," Using the generating function for Hermite polynomials, one can show that, for the natural choice of $\gamma^{2}=\alpha^{2}+\beta^{2}$, the smoothing matrix is given by for $m-n\ge0$ and even $G_{nm}$ vanishes otherwise), and where $\omega^{-2} \equiv \alpha^{-2}+\beta^{-2}$."
+Figure 7 shows how this analytic formula can be used to clliciently smooth a 2-climensional image (see cliscussion in refconvolutionsdbelow)., Figure \ref{fig:smooth} shows how this analytic formula can be used to efficiently smooth a 2-dimensional image (see discussion in \\ref{convolution_2d} below).
+ Anintuilivefeclingfortheef fectof« onvoluliononth for diflerent values of the smoothing scale 3.," An intuitive feeling for the effect of convolution on the shapelet coefficients can be obtained from Figure \ref{fig:gmatrix}, which graphically shows the smoothing matrix $G_{nm}(\alpha,\beta,\gamma=(\alpha^2+\beta^2)^{\frac{1}{2}})$ for different values of the smoothing scale $\beta$ ."
+ As expected. the smoothing matrix approaches the identity matrix in the limitof vanishing smoothing scale (3= 0).," As expected, the smoothing matrix approaches the identity matrix in the limitof vanishing smoothing scale $\beta \rightarrow 0$ )."
+ On the other hand. for verylarge smoothing kernels (30 o5)M reduces to a projection of all the input modes m onto the n=0 output mode.," On the other hand, for verylarge smoothing kernels $\beta \rightarrow \infty$ )it reduces to a projection of all the input modes $m$ onto the $n=0$ output mode."
+ For intermediate scales. the," For intermediate scales, the"
+"As mentioned, the complete elliptic integral K exhibits a logarithmic divergence as its modulus k—1.","As mentioned, the complete elliptic integral $\elik$ exhibits a logarithmic divergence as its modulus $k \rightarrow 1$."
+" This is the direct consequence of the point masssingularity (i.e., |?—|> 0)."," This is the direct consequence of the point masssingularity (i.e., $|\vec{r}-\vec{r'}| \rightarrow 0$ )."
+ The determination of accurate potentials from Eqs., The determination of accurate potentials from Eqs.
+" 3 and 5 is therefore not straightforward and requires a careful treatment of improper integrals (e.g.,Stemwedeletal.1990;Huré2005;Huré&Pierens 2005)."," \ref{eq:psit} and \ref{eq:psif} is therefore not straightforward and requires a careful treatment of improper integrals \citep[e.g.,][]{syc90,hure05,hurepierens05}."
+". This technical difficulty is usually circumvented by changing the relative separation according to where A#0 is a constant known as the ""softening length"".", This technical difficulty is usually circumvented by changing the relative separation according to where $\lambda \ne 0$ is a constant known as the “softening length”.
+ The main drawback is that softened gravity modifies Newton's law for gravitation both on short and long ranges., The main drawback is that softened gravity modifies Newton's law for gravitation both on short and long ranges.
+" It lowers the magnitude of forces, enhances stability, and introduces a bias in models that is not easy to measure and interpret (forstellarRomeo1994;Sommer-Larsenetal.1998;Dehnen 2001)."," It lowers the magnitude of forces, enhances stability, and introduces a bias in models that is not easy to measure and interpret \citep[for stellar and gas discs see, e.g.,][]{paplin89,sayo90,romeo94,sommer98,dehnen01}."
+". In the case of gaseous discs of interest here, the modification of the relative distances according to Eq."," In the case of gaseous discs of interest here, the modification of the relative distances according to Eq."
+ 36 changes the expressions for ΠΕ and Vin., \ref{eq:sl} changes the expressions for $\psif$ and $\psit$ .
+" It is however easy to show that the associated “softened” potentials denoted Vf and Vi, respectively, can still be written in terms of Eqs."," It is however easy to show that the associated “softened” potentials denoted $\psif_{\rm s}$ and $\psit_{\rm s}$, respectively, can still be written in terms of Eqs."
+" 3 and 5,, respectively, provided that the modulus k is replaced by and m is replaced by In disc models and simulations, one never (or rarely) computes Vin, or its fully asymmetric/tri-dimensionalversion (seehoweverLietal.2009)."," \ref{eq:psit} and \ref{eq:psif}, respectively, provided that the modulus $k$ is replaced by and $m$ is replaced by In disc models and simulations, one never (or rarely) computes $\psit$, or its fully asymmetric/tri-dimensionalversion \citep[see however][]{li09}."
+". Instead, one computes Wf*t in the framework of softened gravity,that is Wf, "," Instead, one computes $\psif$ in the framework of softened gravity,that is $\psif_{\rm s}$."
+The softening length A appearing in Eq., The softening length $\lambda$ appearing in Eq.
+ 38 must therefore be prescribed., \ref{eq:msoft} must therefore be prescribed.
+ We note that solving Eq., We note that solving Eq.
+ 38 for A leads to where m/=4/1—m? is the complementary modulus., \ref{eq:msoft} for $\lambda$ leads to where $\ms' = \sqrt{1-\ms^2}$ is the complementary modulus.
+ Many prescriptions for 2 have been proposed., Many prescriptions for $\lambda$ have been proposed.
+" In general, this is not a constant but a certain function of the radius and/or disc parameters."," In general, this is not a constant but a certain function of the radius and/or disc parameters."
+ Table 1 gathers a few formulae for A used by different authors over twenty years., Table \ref{tab:variouspresc} gathers a few formulae for $\lambda$ used by different authors over twenty years.
+" Although not exhaustive, this list clearly shows that there is no trend in magnitude and variation in space (and possible dependency on the disc parameters)."," Although not exhaustive, this list clearly shows that there is no trend in magnitude and variation in space (and possible dependency on the disc parameters)."
+ The results obtained in Sect., The results obtained in Sect.
+ 3 can help substantially to define the appropriate prescription for 4., \ref{sec:chif} can help substantially to define the appropriate prescription for $\lambda$.
+ We can see from Eqs., We can see from Eqs.
+ 5 and 7 that the disc provided that 1 isthe root of the equation for all R., \ref{eq:psif} and \ref{eq:psitbis} that the provided that $\lambda$ isthe root of the equation for all $R$ .
+" Only a numerical approach can yield the exact value of 4, if it exists."," Only a numerical approach can yield the exact value of $\lambda$ , if it exists."
+" However, a good approximation to this root"," However, a good approximation to this root"
+NMuevical siuulatious of hierarchical galaxy formation models siwh as the CDM model predict a niversal cusped ¢ensitv core for the dark iuatter halos of ealaxies {ee. Navarro ct al.,Numerical simulations of hierarchical galaxy formation models such as the CDM model predict a “universal” cusped density core for the dark matter halos of galaxies (e.g. Navarro et al.
+ 1996)., 1996).
+ A cusped deusity core corresponds to a steepvo rising rotation curve., A cusped density core corresponds to a steeply rising rotation curve.
+ The observed snuematics of galaxies cau hence be used to test such nunierical models of galaxy formation., The observed kinematics of galaxies can hence be used to test such numerical models of galaxy formation.
+ Dwart low surface brightuess (LSB). galaxies are best suited for such a fest. slice they. unlike larecr galaxies. are known to ve dark matter donünaut.," Dwarf low surface brightness (LSB), galaxies are best suited for such a test, since they, unlike larger galaxies, are known to be dark matter dominant."
+" In laree spiral galaxies. both eax and stars mnake significant contributions to the total Mass, paricularlv in the iuncr regions of the galaxy."," In large spiral galaxies, both gas and stars make significant contributions to the total mass, particularly in the inner regions of the galaxy."
+ Since he exac contribution of the barvoulc material to the otal mass of the galaxy depeuds on tl known mass o light ratio of the stellar disk. it is generally difficult O Waddyenousily determüne the deusity profile of the dark matter halo in the ceutral regions of large spirals.," Since the exact contribution of the baryonic material to the total mass of the galaxy depends on the unknown mass to light ratio of the stellar disk, it is generally difficult to unambiguously determine the density profile of the dark matter halo in the central regions of large spirals."
+ Iu dwarf LSB galaxies on the other haud. since the stellar disk dis exsnerally cvnamically unmuüportaut. the central halo deusitv can be munch better constrained.," In dwarf LSB galaxies on the other hand, since the stellar disk is generally dynamically unimportant, the central halo density can be much better constrained."
+ Iuterestiugly. the observed rotation curves of dwarf galaxies generally ndicate that their dark iuatter halos have coustaut density cores (e.g. Weldrake et al.," Interestingly, the observed rotation curves of dwarf galaxies generally indicate that their dark matter halos have constant density cores (e.g. Weldrake et al."
+ 2003) uulike the cusped density cores predicted iu uunuerical simulations., 2003) unlike the cusped density cores predicted in numerical simulations.
+ Another prediction of the ποτΊσα. simulations is that the density of the dark matter halo is related to the backerouud deusitv at the time of the halo formation., Another prediction of the numerical simulations is that the density of the dark matter halo is related to the background density at the time of the halo formation.
+" Since the smallest ealaxics fori first iu such modes, these galaxies aro σος‘ted to have the lareest halo densities."," Since the smallest galaxies form first in such models, these galaxies are expected to have the largest halo densities."
+" The deteriunuation of the shapes aud characeristic densities of the dark matter halos of the faintes dwarf galaxies is hence a particularly interesting probαμ,", The determination of the shapes and characteristic densities of the dark matter halos of the faintest dwarf galaxies is hence a particularly interesting problem.
+ However. a nuajor stiuublius block iu such prograus ds that it is currently controversal whether verv fain dwarf nreeular ealaxies show systematic rotation or not.," However, a major stumbling block in such programs is that it is currently controversial whether very faint dwarf irregular galaxies show systematic rotation or not."
+ Lo et al. (, Lo et al. (
+1993). iu a stilv of the kinematics a saie of nine faint dwarfs (xith Mp~9.0 to Mp~ 11.0) found that most of them were characterized by chaotic velocity fields.,"1993), in a study of the kinematics a sample of nine faint dwarfs (with $_B\sim-9.0$ to $_B\sim-14.0$ ) found that most of them were characterized by chaotic velocity fields."
+" Πωπονα, as pointed out bv Skillman(1996)... Lo et al"," However, as pointed out by \cite{skillman96}, Lo et al."
+/s observations had limited sensitivity to faut extended cussion which is likely to have led to an nuderestimation of the rotation velocities.,'s observations had limited sensitivity to faint extended emission which is likely to have led to an underestimation of the rotation velocities.
+ Further. from. a recent hieh seuxitivitv aud Ligh velocity resolution study of the dwarf meegular galaxy. Camelopardalis D (Cin D). Beemctal.(2003) found that inspite of being very faint (Mp~ 10.5) the ealaxy shows systematic rotation.," Further, from a recent high sensitivity and high velocity resolution study of the dwarf irregular galaxy, Camelopardalis B (Cam B), \cite{begum03} found that inspite of being very faint $M_B\sim -10.8$ ), the galaxy shows systematic rotation."
+ Do all faint dwarf regular ealaxics lave a rotating III disk or is Cu D a special case?, Do all faint dwarf irregular galaxies have a rotating HI disk or is Cam B a special case?
+ What is the dark matter distribution iu these very. faint galaxies?, What is the dark matter distribution in these very faint galaxies?
+ Iu this paper we discuss these questious iu the specific context of the local group dwarf galaxy DDO210., In this paper we discuss these questions in the specific context of the local group dwarf galaxy DDO210.
+ DDO210. the faintest known (Mp~—10.6) gas vich dwarf galaxy in our local group. was discovered by van deu Bere (1959) aud later detected in an II 21 cn survey by Fisher aud Tully (1975).," 	DDO210, the faintest known $M_B\sim-10.6$ ) gas rich dwarf galaxy in our local group, was discovered by van den Berg (1959) and later detected in an HI 21 cm survey by Fisher and Tully (1975)."
+ Fisher aud Tully (1979) assigned a distance of 0.7 Alpe to it. based on its proximity to NGC 6822 both on the sky aud iu velocity.," Fisher and Tully (1979) assigned a distance of 0.7 Mpc to it, based on its proximity to NGC 6822 both on the sky and in velocity."
+ Ou the other haud. Caceeioetal.(1993).. based ou the colourauagnitude (C-AL) diagram «© DDO210. estimated its distance to be 2.5 AIpe.," On the other hand, \cite{greggio93}, based on the colour-magnitude (C-M) diagram of DDO210, estimated its distance to be 2.5 Mpc."
+ However. receut distance estimates for DDO210 eive distances closer to the original estimate of Fisher aud Tully (1979). Leeetal.(1999)," However, recent distance estimates for DDO210 give distances closer to the original estimate of Fisher and Tully (1979). \cite{lee99},"
+... based on the I magnitude of the tip of the red giaut. brauch. estimated the distance to DD210 to be 950450 kpe.," based on the I magnitude of the tip of the red giant branch, estimated the distance to DD210 to be $\pm$ 50 kpc."
+ This estimate is in excelleut agreement with the value of LO LO kpe derived receutly by IEaracheutsevetal.(2002) using UST observations., This estimate is in excellent agreement with the value of $\pm$ 40 kpc derived recently by \cite{karachentesv02} using HST observations.
+ At this distauce. DDO210 would be a member of the loca eroup.," At this distance, DDO210 would be a member of the local group."
+" DDO210 is classified as a dblr/dSph or ""trausition ealaxv. with properties intermediate between dwarf regulus and dwarf spheroidals (Mateo 1998)."," DDO210 is classified as a dIr/dSph or “transition galaxy”, with properties intermediate between dwarf irregulars and dwarf spheroidals (Mateo 1998)."
+ For exinuple. in spite of containing a significant amount of neutral gas. DDO210 shows no sigus of ongoing star formation.," For example, in spite of containing a significant amount of neutral gas, DDO210 shows no signs of ongoing star formation."
+ Ta imaging detected a single source of line cuuission in the galaxy: however follow up observations of this cuuission sugecsts that it docs not ariso m a norma III region. but probably comes from deuse outflowiusg material from an evolved star (wan Zee ct al.," $\alpha$ imaging detected a single source of line emission in the galaxy; however follow up observations of this emission suggests that it does not arise in a normal HII region, but probably comes from dense outflowing material from an evolved star (van Zee et al."
+ 1997)., 1997).
+produced [ie.. 9aa(n.£D)—02.(n.D) p.(n.D)].,"produced [i.e., $\delta\nu_{\mathrm {model}}({\it n,l})-\delta\nu_\odot({\it n,l})$ $\nu_\odot({\it n,l})$ ]."
+ As in the case of the frequency difference plots. the closer 0aqan.£)—05.(n.E) is to 0 jllz. the better the agreement between the observations and a mocel.," As in the case of the frequency difference plots, the closer $\delta\nu_{\mathrm {model}}({\it n,l})-\delta\nu_\odot({\it n,l})$ is to 0 $\mu$ Hz, the better the agreement between the observations and a model."
+ since the sensitivity of the small spacines to the deep interior of (he Sun diminishes with increasing4 only those panodes with 0. 1. and 2 were used to contrast the models with observed values Grom the Sun.," Since the sensitivity of the small spacings to the deep interior of the Sun diminishes with increasing, only those -modes with =0, 1, and 2 were used to contrast the models with observed values from the Sun."
+ The small spacing difference plots can be seen in Figure 3., The small spacing difference plots can be seen in Figure 3.
+" Figure3 is a 3x5 grid in whic| each row contains moclels οἱ the same Zin. where Zi Zinii includes models #11-4. Z;,,—0.50Z;,4"" includes models 55-8. Z;i,,70.65Z;,4, includes models i:09-12. Zi - ine models 3£113-16. and Zing =Zing includes the standard solar models 4117-20."," Figure 3 is a $\times$ 5 grid in which each row contains models of the same $_{\mathrm {int}}$, where $_{\mathrm {int}}$ $_{\mathrm {init}}$ includes models 1-4, $_{\mathrm {int}}$ $_{\mathrm {init}}$ includes models 5-8, $_{\mathrm {int}}$ $_{\mathrm {init}}$ includes models 9-12, $_{\mathrm {int}}$ $_{\mathrm {init}}$ includes models 13-16, and $_{\mathrm {int}}$ $_{\mathrm {init}}$ includes the standard solar models 17-20."
+ Each column in Figure 3 contains results for a common Evalue. wilh 0 in column 1. 21 in column 2. and /—2 in column 3.," Each column in Figure 3 contains results for a common -value, with =0 in column 1, =1 in column 2, and =2 in column 3."
+ Lines in each panel represent models with (he same Ziy with differing μμ values ranging from 0.0170 to 0.0220. as indicated.," Lines in each panel represent models with the same $_{\mathrm {int}}$ with differing $_{\mathrm {init}}$ values ranging from 0.0170 to 0.0220, as indicated."
+ Figure 4 shows the relative souud-speed differences lor the twenty. models presented in (his paper compared to the observed solar sound-speed. derived [rom an inversion of the solar p-mode frequencies (BPBOO).," Figure 4 shows the relative sound-speed differences for the twenty models presented in this paper compared to the observed solar sound-speed, derived from an inversion of the solar -mode frequencies (BPB00)."
+ We note the superior agreement of the standard solar models with observation. especially model #220. as compared to the non-standard mocels.," We note the superior agreement of the standard solar models with observation, especially model 20, as compared to the non-standard models."
+ Sinularly. Figure5 shows the same plots for densitv.," Similarly, Figure 5 shows the same plots for density."
+ The less precise agreement with densities (han with sound-speedis in part a reflection of the greater uncertainty in clensily inversions., The less precise agreement with densities than with sound-speed is in part a reflection of the greater uncertainty in density inversions.
+ Figure 1 shows the characteristics of the standard solar models (2:117-20) in the frequency difference diagram., Figure 1 shows the characteristics of the standard solar models 17-20) in the frequency difference diagram.
+ Best agreement with observation. as measured bv the thickness of the line bundles corresponding to dillerent /-values. favors #220 and 19.," Best agreement with observation, as measured by the thickness of the line bundles corresponding to different $l$ -values, favors 20 and 19."
+ Note that our model #119 is nearly identical to model 4220 of GD97. (hal was judged by GD9T as (he “best” stanclarel solar model in their study.," Note that our model 19 is nearly identical to model 20 of GD97, that was judged by GD97 as the “best” standard solar model in their study."
+" But our model 4220."" which is intermediate in Zi between GDO9'Ts models 4220 and 21. is a better moclel"," But our model 20, which is intermediate in $_{\mathrm {init}}$ between GD97's models 20 and 21, is a better model still."
+ The bundle of /values in model #220 is thinner than in model 4119. a fact which is also reflected in (he best agreement with observation in (le souid-speedplot shown in Figure 4.," The bundle of $l$ -values in model 20 is thinner than in model 19, a fact which is also reflected in the best agreement with observation in the sound-speedplot shown in Figure 4."
+ In selecting a “best” standard solar model. one should also take into account two," In selecting a “best” standard solar model, one should also take into account two"
+Our current understanding of galaxy formation has greatly venehted from the results of ΑΝ body iodelling of structure formation in the carly Universe. which predicts dat small objects combine eravitationally to produce je ealaxies we see today: a xocess called. hicrarchical-chisterine-ereine (hereafter HCNL cf.,"Our current understanding of galaxy formation has greatly benefited from the results of $N-$ body modelling of structure formation in the early Universe, which predicts that small objects combine gravitationally to produce the galaxies we see today: a process called hierarchical-clustering-merging (hereafter HCM, cf."
+ Kauffinaunn ct , Kauffmann et al.
+1993)., 1993).
+ One of the tenets of the ΠΟΝΤ paraciei is iat galaxies are constantly mereing with one another., One of the tenets of the HCM paradigm is that galaxies are constantly merging with one another.
+ Iu the case of elliptical aud SO ealaxies. there ds Huple observational evidence that they are contiuuallv subjected to mergers with snaller. neiglibouriug galaxies (c£.," In the case of elliptical and S0 galaxies, there is ample observational evidence that they are continually subjected to mergers with smaller, neighbouring galaxies (cf."
+ Sclsveizer 1998)., Schweizer 1998).
+ If the ICAL paradigni is uuiversal. spirals are subject to the same formation processes as E aud SOs.," If the HCM paradigm is universal, spirals are subject to the same formation processes as E's and S0's."
+ Often the fugerpriuts of such secoud events reside in the stellar and/or gaseous kinematics of a ealaxy rather than iu its morphology., Often the fingerprints of such second events reside in the stellar and/or gaseous kinematics of a galaxy rather than in its morphology.
+ This is particularly true if we consider that the most evident ‘morphological tracers” of interactions such as peculiar or spindle galaxies make up less than 5% of all objects in azuv one of the RC3 (de Vaucouleurs et al., This is particularly true if we consider that the most evident `morphological tracers' of interactions such as peculiar or spindle galaxies make up less than $5\%$ of all objects in any one of the RC3 (de Vaucouleurs et al.
+ 1991). UGC (Nilsou. 1973) or ESO/Upssala (Lauberts 1982) galaxw catalogues.," 1991), UGC (Nilson 1973) or ESO/Upssala (Lauberts 1982) galaxy catalogues."
+ It is therefore crucial to obtain detailed kincmatic parameters of both stars and eas to unveil the relics of accretion or mereing eveuts which have occurred in galaxy history., It is therefore crucial to obtain detailed kinematic parameters of both stars and gas to unveil the relics of accretion or merging events which have occurred in galaxy history.
+ A large fraction of spirals exhibit kinematic disturbances raneine from iud to major. and can ecuerally be explaiue as the visible sigus of tidal encounters (Rubin. Waterman Kenney 1999).," A large fraction of spirals exhibit kinematic disturbances ranging from mild to major, and can generally be explained as the visible signs of tidal encounters (Rubin, Waterman Kenney 1999)."
+" In recent vears a uuuber of otherwise unorphologically uudisturbed spirals have becu found which host |inematically-decoupled components (ISDC""«). such as stellar KDC's (Bertola et al."," In recent years a number of otherwise morphologically undisturbed spirals have been found which host kinematically-decoupled components (KDC's), such as stellar KDC's (Bertola et al."
+ 1999: Sarzi ct al., 1999; Sarzi et al.
+ 2000). counter-rotating exteuded stellar discs (Alerrifield I&uijkeu 1901: Bertola et al.," 2000), counter-rotating extended stellar discs (Merrifield Kuijken 1994; Bertola et al."
+ 1996: Jore et al., 1996; Jore et al.
+ 1996). counter-rotatiug or decoupled gaseous disces (Brawn et al.," 1996), counter-rotating or decoupled gaseous discs (Braun et al."
+ 1992: Rubin 1991: Ris et al., 1992; Rubin 1994; Rix et al.
+ 1995: Cir et al., 1995; Ciri et al.
+ 1995: Tavues et al., 1995; Haynes et al.
+ 2000: Kannappan Fabricant 2001) and possibly couuter-rotatiug lees (Prada et al., 2000; Kannappan Fabricant 2001) and possibly counter-rotating bulges (Prada et al.
+ 1996: but see also Bottema 1999)., 1996; but see also Bottema 1999).
+ Studving the interplay between ionized eas and stellar kinematics allows us to address other issues concerning the dyainical structure of spirals., Studying the interplay between ionized gas and stellar kinematics allows us to address other issues concerning the dynamical structure of spirals.
+ These iuclude the origiu of disc heating and the presence of stellar or gaseous discs in galactic nuclei., These include the origin of disc heating and the presence of stellar or gaseous discs in galactic nuclei.
+ Caavitational scattering from eiaut molecular clouds aux spiral density waves are tle prime candidates to explain the finite thickness of stellar discs., Gravitational scattering from giant molecular clouds and spiral density waves are the prime candidates to explain the finite thickness of stellar discs.
+ It is expected that the onunant beating mechanisiu varies along the IIubble sequence nit up to now oulv two external galaxies have been studied iu detail (Cterssen. Iujken Merrifield 1996. 2000).," It is expected that the dominant heating mechanism varies along the Hubble sequence but up to now only two external galaxies have been studied in detail (Gerssen, Kuijken Merrifield 1996, 2000)."
+ The presence in the nuclei of SOs aud spirals of small stellar (Eaisellem et al., The presence in the nuclei of S0's and spirals of small stellar (Emsellem et al.
+ 1996: Iormendy ct al., 1996; Kormendy et al.
+" 199G6a.b: van deu Bosch. Jaffe van der Marel 1998: Scorza van deu Bosch 1998: van den Bosch Euiselleii 1998) and/or gaseous discs (Rubin. Ποιον, Young 1997: Bertola ct al."," 1996a,b; van den Bosch, Jaffe van der Marel 1998; Scorza van den Bosch 1998; van den Bosch Emsellem 1998) and/or gaseous discs (Rubin, Kenney, Young 1997; Bertola et al."
+ 1998: Funes 2000) is usually connected to the presence of a ceutral mass concentration., 1998; Funes 2000) is usually connected to the presence of a central mass concentration.
+ It also appears that the central black-hole mass is very stronely correlated with the stellar velocity dispersion of the host galaxwv bulee as recently found bv different authors (Ferrarese \lerritt 2000: Gebhardt et al., It also appears that the central black-hole mass is very strongly correlated with the stellar velocity dispersion of the host galaxy bulge as recently found by different authors (Ferrarese Merritt 2000; Gebhardt et al.
+ 2000)., 2000).
+ This relation is however used on samples which are affected bv. different biases aud therefore new black-hole, This relation is however based on samples which are affected by different biases and therefore new black-hole
+"The mainproduction of leading X. =""and ti comes [rom (heconverted. spin-1","branching ratio of $\Sigma^*\rightarrow\Sigma\pi$ is $\%$, the effect of $\Sigma^*$ decay on $\Sigma^-$ polarization is small."
+This investigationBS was supported in part by NASA 5grant NNGOSGASIG. We have mace use of the data products from the Two Micron All Sky Survey. which is a joint project of the University of Massachusetts and the Infrared Processing aud Analysis Center/Califoraia Institute of Technology. funded by the National Aeronauties and Space Administration and the National Science Foundation.,"This investigation was supported in part by NASA grant NNG05GA84G. We have made use of the data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."
+ We have also used the USNOFS Image and Catalog Archive, We have also used the USNOFS Image and Catalog Archive
+his rule. in that BATSE trigger 2665 does not appear to iive à power law relation between 7 and ££.,"this rule, in that BATSE trigger 2665 does not appear to have a power law relation between $\tau$ and $E$."
+ Furthermore. the exponent of the power Ίαν relationship between £ and τ does not appear to be consistent between GRBs.," Furthermore, the exponent of the power law relationship between $E$ and $\tau$ does not appear to be consistent between GRBs."
+" For example. the dominant. pulse in BATSE trigger 2197 appears well fit by rox4""7. vielding a relatively large difference between the duration of the pulse at high and low eneries."," For example, the dominant pulse in BATSE trigger 2197 appears well fit by $\tau \propto E^{-0.5}$, yielding a relatively large difference between the duration of the pulse at high and low eneries."
+" In contrast. the dominant pulse in BATSE trigger 7087 appears well fit by zx££IT, vielding a relatively small dillerence between the duration of the pulse at high and low energies."," In contrast, the dominant pulse in BATSE trigger 7087 appears well fit by $\tau \propto E^{-0.17}$, yielding a relatively small difference between the duration of the pulse at high and low energies."
+ Note that the dillerence between the 7 [actors of two energy channels is conceptually similar to the lag between hese two energy channels., Note that the difference between the $\tau$ factors of two energy channels is conceptually similar to the lag between these two energy channels.
+ Lt is therefore possible that the power law exponent between 7 and ££ might act similar to )ulse lag and be a measure of intrinsic brightness of the ;ulse. and hence the CRB., It is therefore possible that the power law exponent between $\tau$ and $E$ might act similar to pulse lag and be a measure of intrinsic brightness of the pulse and hence the GRB.
+ Since 7r is the time between 1e pulse. peak and the pulse start. the dillerence between wo rs of the same pulse at dilferent. energies is just. the ference between the peak times of the pulse at. those snereies.," Since $\tau$ is the time between the pulse peak and the pulse start, the difference between two $\tau$ s of the same pulse at different energies is just the difference between the peak times of the pulse at those energies."
+ Although the lag is formally computed by noting 1e time of the maximum in the cross correlation between 1e light curves of two energy channels. this cross correlation would have 1Us greatest instantaneous power when the two »vaks are aligned.," Although the lag is formally computed by noting the time of the maximum in the cross correlation between the light curves of two energy channels, this cross correlation would have it's greatest instantaneous power when the two peaks are aligned."
+ Therefore lag is strongly pulled. toward 1e dillerence in peak times., Therefore lag is strongly pulled toward the difference in peak times.
+ Therefore. one would expect aga.~(Tsτι).," Therefore, one would expect ${\rm lag}_{23} \sim (\tau_3 - \tau_1)$."
+ tis evident [rom inspection that CRB pulses have similarities at dillerent energies., It is evident from inspection that GRB pulses have similarities at different energies.
+ Llere clear rules are postulated. for. transferring specific. pieces of information between energies., Here clear rules are postulated for transferring specific pieces of information between energies.
+" Specific analyses presented above indicate that at least one class of GRB pulses. here dubbed. ""energy scalable pulses” (ESPs). can be described by a simple."," Specific analyses presented above indicate that at least one class of GRB pulses, here dubbed “energy scalable pulses"" (ESPs), can be described by a simple,"
+"Seven objects within the sample have been classified as optical AGNs from previous surveys (see GA09 and Table 1) using typical optical emission-line diagnostics (e.g., the Baldwin-Phillips-Terlevich diagnostic diagrams; ? 1981); however, all have detected emission, and thus are unambiguously identified to host AGNs at mid-IR wavelengths We find that with the exception of NGC 5128 (Centaurus A), our sample is dominated by AGNs with SMBHs in the mass range Mpy&(0.1-- 5)x10'Mc (median of Mguz7x106 Mc).","Seven objects within the sample have been classified as optical AGNs from previous surveys (see GA09 and Table 1) using typical optical emission-line diagnostics (e.g., the Baldwin-Phillips-Terlevich diagnostic diagrams; \citeauthor{bpt} 1981); however, all have detected emission, and thus are unambiguously identified to host AGNs at mid-IR wavelengths We find that with the exception of NGC 5128 (Centaurus A), our sample is dominated by AGNs with SMBHs in the mass range $\Mbh
+\approx (0.1$ $5) \times 10^7 \Msun$ (median of $\Mbh \approx 7
+\times 10^6 \Msun$ )."
+" Due to the irregular structure of one of the galaxies in the sample (NGC 5195), Mau is poorly determined; in Fig."," Due to the irregular structure of one of the galaxies in the sample (NGC 5195), $\Mbh$ is poorly determined; in Fig."
+" 3 we plot Mgu estimates from both the M-o, and Msan-Lx,»u relations (connected blue-dashed line)."," \ref{fig_2} we plot $\Mbh$ estimates from both the $\sigma_*$ and $\Mbh$ $_{\rm K,bul}$ relations (connected blue-dashed line)."
+" We find the AGNs in our sample are spread over a wide range of bolometric luminosities, Lgoi.acn1010— 1075ergs!."," We find the AGNs in our sample are spread over a wide range of bolometric luminosities, $L_{\rm Bol,AGN} \approx 10^{40}$ $10^{45}
+\ergps$."
+" To assess the relative mass-accretion rates of the sample (LgoiAGN/Lmgaa~ 7), we over-plot lines of constant Eddington ratios (η&107~3,107*, derived following ?)) and their associated mass-doubling 1.0;times (tz30,0.3,0.03 Gyrs, respectively)."," To assess the relative mass-accretion rates of the sample $L_{\rm Bol,AGN} / L_{\rm Edd} \sim \eta$ ), we over-plot lines of constant Eddington ratios $\eta \approx 10^{-3},10^{-1},1.0$; derived following \citealt{rees84}) ) and their associated mass-doubling times $t \approx 30,0.3,0.03$ Gyrs, respectively)."
+" Given the large range in bolometric luminosities, it is not surprising that the AGNs in the sample are found to be accreting at rates covering over 5 orders of magnitude (ηzz10 ?—1)."," Given the large range in bolometric luminosities, it is not surprising that the AGNs in the sample are found to be accreting at rates covering over 5 orders of magnitude $\eta \approx 10^{-5}$ –1)."
+" With the exception of a few AGNs, the observed range in Eddington ratios is found to be roughly consistent with those found by H04 for active galaxies (solid contours in Fig. 3))."," With the exception of a few AGNs, the observed range in Eddington ratios is found to be roughly consistent with those found by H04 for active galaxies (solid contours in Fig. \ref{fig_2}) )."
+" As our work is not limited by the spectral resolution of the SDSS (i.e., with a limit of Msn>3x109Mo), we show in Fig."," As our work is not limited by the spectral resolution of the SDSS (i.e., with a limit of $\Mbh \ga 3 \times 10^6 \Msun$ ), we show in Fig."
+" 3 that significant accretion, η>107° (i.e., radiatively efficient accretion systems; e.g., thin discs) occurs onto SMBHs with Mgg£(1-3) x105Mo."," \ref{fig_2} that significant accretion, $\eta > 10^{-3}$ (i.e., radiatively efficient accretion systems; e.g., thin discs) occurs onto SMBHs with $\Mbh \approx (1$ $3) \times 10^6 \Msun$."
+" The majority of these low-mass, rapidly-accreting SMBHs are hosted in late-type, disc-dominated spiral galaxies (Sc-Sd)."," The majority of these low-mass, rapidly-accreting SMBHs are hosted in late-type, disc-dominated spiral galaxies (Sc–Sd)."
+" By contrast, it is generally assumed that gas-rich late-type spirals are preferentially inactive galaxies and that a large bulge may be a necessary component for the existence of a SMBH, and thus a luminous AGN."," By contrast, it is generally assumed that gas-rich late-type spirals are preferentially inactive galaxies and that a large bulge may be a necessary component for the existence of a SMBH, and thus a luminous AGN."
+" Furthermore, of the four AGNs within the sample with SMBHs consistent with Mguz106Mo, we find that three sources are not identified as AGNs in sensitive optical surveys."," Furthermore, of the four AGNs within the sample with SMBHs consistent with $\Mbh \approx 10^6
+\Msun$, we find that three sources are not identified as AGNs in sensitive optical surveys."
+ This indicates that significant SMBH accretion may be missed by statistically-large optical surveys such as H04 even if the spectral resolution was sufficient to identify SMBHs down to 10° Mo., This indicates that significant SMBH accretion may be missed by statistically-large optical surveys such as H04 even if the spectral resolution was sufficient to identify SMBHs down to $\Mbh \approx 10^6 \Msun$ .
+" For the subset of our AGN sample which host SMBHs with MauZ3x10°Mo, we find that many of the optically unidentified AGNs are accreting at relatively low Eddington ratios (η$10?), and are unlikely to make a significant additional contribution to the present-day growth of SMBHs."," For the subset of our AGN sample which host SMBHs with $\Mbh \goa 3
+\times 10^6 \Msun$, we find that many of the optically unidentified AGNs are accreting at relatively low Eddington ratios $\eta \loa
+10^{-3}$ ), and are unlikely to make a significant additional contribution to the present-day growth of SMBHs."
+" However, these same AGNs may form part of a separate, underlying population of radiatively inefficient accretion systems such as advection dominated accretion flows (ADAFs; e.g., ?)) or those which contain optically-thick slim-discs."," However, these same AGNs may form part of a separate, underlying population of radiatively inefficient accretion systems such as advection dominated accretion flows (ADAFs; e.g., \citealt{adaf}) ) or those which contain optically-thick slim-discs."
+" Further spectral analysis of the X-ray data may distinguish between these particular accretion systems, but is beyond the scope of these analyses (see Goulding et al."," Further spectral analysis of the X-ray data may distinguish between these particular accretion systems, but is beyond the scope of these analyses (see Goulding et al."
+ in preparation)., in preparation).
+" Using the relative mass accretion rates estimated for our sample (Fig. 3)),"," Using the relative mass accretion rates estimated for our sample (Fig. \ref{fig_2}) ),"
+ we can infer the volume-averaged growth time of SMBHs in the local Universe., we can infer the volume-averaged growth time of SMBHs in the local Universe.
+" Assuming a mean Kerr spin parameter (a) for our sample of az0.67 (e.g., ??)), Le., an accretion efficiency (e) of z0.1, the characteristic mass doubling time (ον) of a SMBH accreting matter at the Eddington limit is tea;+30 Myrs (Rees 1984)."," Assuming a mean Kerr spin parameter $a$ ) for our sample of $a \approx 0.67$ (e.g., \citealt{treister06,hopkins07}) ), i.e., an accretion efficiency $\epsilon$ ) of $\approx 0.1$, the characteristic mass doubling time $t_{2M}$ ) of a SMBH accreting matter at the Eddington limit is $t_{2M} \approx 30$ Myrs (Rees 1984)."
+" Under the further assumption that a, and hence ε, does not vary significantly for changes in Mpu (?),, we assess the present-day growth rate of Following H04, we calculate and extend to lower masses (Mau<3x109 Ma) the integrated growth of SMBHs."," Under the further assumption that $a$, and hence $\epsilon$, does not vary significantly for changes in $\Mbh$ \citep{king08}, we assess the present-day growth rate of Following H04, we calculate and extend to lower masses $\Mbh < 3 \times 10^6
+\Msun$ ) the integrated growth of SMBHs."
+ Growth time errors are calculated from the log-normal standard deviations of the sample., Growth time errors are calculated from the log-normal standard deviations of the sample.
+ We note here that we also include the optically unidentified AGNs which would not be detected in the SDSS., We note here that we also include the optically unidentified AGNs which would not be detected in the SDSS.
+" In Fig. 4,,"," In Fig. \ref{fig_4},"
+" we find that the mean growth time for low- SMBHs (Mauz&10° Mo) is z6*9 Gyrs, which is consistent with these AGNs growing on time-scales similar to that of the age of the Universe."," we find that the mean growth time for low-mass SMBHs $\Mbh \approx 10^6 \Msun$ ) is $\approx 6^{+6}_{-3}$ Gyrs, which is consistent with these AGNs growing on time-scales similar to that of the age of the Universe."
+ Our results are found to be broadly consistent with a simple extrapolation of the, Our results are found to be broadly consistent with a simple extrapolation of the
+with oA15 can fit the angular cdiameter distance to decoupling measured by WNLAPT. so these are of particular interest as hey are consistent with current observations of the background. but may on the other hand allect the erowth of structures.,"with $\beta_0 \leq 0.15$ can fit the angular diameter distance to decoupling measured by WMAP7, so these are of particular interest as they are consistent with current observations of the background, but may on the other hand affect the growth of structures."
+" 1t can be seen that there is a difference in the Ξ0 powerspectrum between ACDAL and EX""001. the lowest of the couplings investigated here. and a dillerence between ACDAL and the highest of the COMings. EXPOOS."," It can be seen that there is a difference in the $z=0$ powerspectrum between $\Lambda$ CDM and EXP001, the lowest of the couplings investigated here, and a difference between $\Lambda $ CDM and the highest of the couplings, EXP003."
+ In this analysis we do not use the simulated: power spectrum directly. but instead. usethe ratio between the ACDM and eDIS power spectra to find the dillerence in the erowth of modes for dillerent couplings with the same initial conditions., In this analysis we do not use the simulated power spectrum directly but instead usethe ratio between the $\Lambda$ CDM and cDE power spectra to find the difference in the growth of modes for different couplings with the same initial conditions.
+ Using this method reduces the error associated with the limited number of independent A-modes that enter the computation of the power in each & bin to only the error on the ACDAL power spectrum., Using this method reduces the error associated with the limited number of independent $k$ -modes that enter the computation of the power in each $k$ bin to only the error on the $\Lambda$ CDM power spectrum.
+ Now we present the framework for calculating the &ravitational lensinge signalo in the eDIZ scenario., Now we present the framework for calculating the gravitational lensing signal in the cDE scenario.
+ The wav that light is dellected along the path from its source to an observer is determined. by the mass distribution and. the ecometry of the Universe., The way that light is deflected along the path from its source to an observer is determined by the mass distribution and the geometry of the Universe.
+ The dellections of light lead. to distortions of the observed image of the source., The deflections of light lead to distortions of the observed image of the source.
+ The mapping between the original source shape and the observed. image is given by (Bartelmann&Schneider2001) where the convergence. H. Causes an isotropic dilation and the shear. τσι|752. changes the ellipticitv.," The mapping between the original source shape and the observed image is given by \citep[][]{Bartelmann:1999yn} where the convergence, $\kappa$, causes an isotropic dilation and the shear, $\gamma=\gamma_1+i\gamma_2$, changes the ellipticity."
+ & is challengingto measure. as the original size of the source is unknown: equally ~ cannot oe measured for a single source as the intrinsic ellipticity of the source is unknown.," $\kappa$ is challengingto measure, as the original size of the source is unknown; equally $\gamma$ cannot be measured for a single source as the intrinsic ellipticity of the source is unknown."
+ However if the shear of a Large number of sources is correlated. then the lensing signal can »e measured as a correlation function (insofar as the intrinsic ellipticities are not. themselves correlated: see. discussion in section 7? below).," However if the shear of a large number of sources is correlated, then the lensing signal can be measured as a correlation function (insofar as the intrinsic ellipticities are not themselves correlated; see discussion in section \ref{Results} below)."
+" Therefore we will be interested. in 16 shear correlation function C- in order to quantify our oedictions. given by (Bartelmann&Schneider2001) where 8 is the angular distance between the correlated sources. / is the angular wavenumber and the lensing power spectrum 27, is given by (Baconetal.2005:2007) with the weight functions where X is comoving distance. yg is the comoving distance to the horizon and Cy) is the normalised distribution of the SOULCOS in Comoving distance. corresponding to a redshift) distribution for the sources."," Therefore we will be interested in the shear correlation function $C_\gamma$ in order to quantify our predictions, given by \citep[][]{Bartelmann:1999yn}
+ where $\theta$ is the angular distance between the correlated sources, $l$ is the angular wavenumber and the lensing power spectrum $P_\kappa$ is given by \citep[][]{Bacon:2004ht,Massey:2007gh}
+ with the weight functions where $\chi$ is comoving distance, $\chi _{\rm H}$ is the comoving distance to the horizon and $G(\chi)$ is the normalised distribution of the sources in comoving distance, corresponding to a redshift distribution for the sources."
+ We use two weight5 functions in Equation1 14. since we are using tomographic weak lensing., We use two weight functions in Equation \ref{eq:Pkappa} since we are using tomographic weak lensing.
+ Equation (15)) is valid for lat cosmologies.5 which are all that are considered. in this yaper.," Equation \ref{eq:w}) ) is valid for flat cosmologies, which are all that are considered in this paper."
+ Usually Eq. (14)), Usually Eq. \ref{eq:Pkappa}) )
+" is written. with the assumption Quail)=Qu. 0: however the form above does not include such an assumption. as coupling CDM. and DE means hat Q,, has a different dependence on time. as shown in ]5qs. (1--4))."," is written with the assumption $\Omega_{\rm m}(a)=\Omega_{\rm m}/a^3$ ; however the form above does not include such an assumption, as coupling CDM and DE means that $\Omega_{\rm m}$ has a different dependence on time, as shown in Eqs. \ref{klein_gordon}- \ref{continuity_radiation}) )."
+ We have modified the COSMOS. CosmoMC. code (Lesgourguesetal.2007:Lewis&Bricle2002:Masseyal. 2007).. which calculates the combined shear correlation unction from the theoretical power spectrum. prediction eiven by CosmoMC. to include eross-correlation of recshift jns and to calculate. the predicted. weak lensing signal directly. from the eDIZ model power spectra. according to σαν. (13--15)).," We have modified the COSMOS CosmoMC code \citep[][]{Lesgourgues:2007te,Lewis:2002ah,Massey:2007gh}, which calculates the combined shear correlation function from the theoretical power spectrum prediction given by CosmoMC, to include cross-correlation of redshift bins and to calculate the predicted weak lensing signal directly from the cDE model power spectra, according to Eqs. \ref{eq:ckappa}- \ref{eq:w}) )."
+ Wewill now use these results to estimate the discriminatory power from lensing between cillerent coupled DE mocdels., Wewill now use these results to estimate the discriminatory power from lensing between different coupled DE models.
+ For our analysis we will rely on he public nonlinear power spectrum data computed from the simulations (Baldi2011b).. the largest suite of cosmological N-bocly simulations for cDI moclels to date. carried out with the moclified version by Daldietal.(2010) of the widely used Jrec-P parallel N-body code. (Springel 2005)..," For our analysis we will rely on the public nonlinear power spectrum data computed from the simulations \citep[][]{CoDECS}, the largest suite of cosmological N-body simulations for cDE models to date, carried out with the modified version by \citet{Baldi_etal_2010} of the widely used Tree-PM parallel N-body code \citep[][]{gadget-2}. ."
+ In particularM. we will consider the suite that includes hverodvnamical simulations o Pall the EeDE models summarized in Table 1. on relativelysmall scales., In particular we will consider the suite that includes hydrodynamical simulations of all the cDE models summarized in Table \ref{tab:models} on relativelysmall scales.
+ More specifically. the runs follow he evolution of 512° ," More specifically, the runs follow the evolution of $512^{3}$ "
+[16]..,\cite{alan}.
+" The noncommutative spaces can be realized as spaces where coordinateoperators &, satisfy the commutation relations: [2,,ἂν]=ἶθμν where θµν is a real antisymmetric tensor with the dimension of [L]?."," The noncommutative spaces can be realized as spaces where coordinateoperators $\hat{x}_{\mu}$ satisfy the commutation relations: $[\hat{x}_{\mu},\hat{x}_{\nu}]=i\theta_{\mu\nu}$ where $\theta_{\mu\nu}$ is a real antisymmetric tensor with the dimension of $[L]^{2}$."
+" We note that a space-time noncommutativity, (ιο4 0), might lead to some problems with unitarity and causality."," We note that a space-time noncommutativity, $\theta_{i0} \neq 0$ ), might lead to some problems with unitarity and causality."
+" In [0],, we have calculated the dependence of the charged lepton spin analyzing power in the top quark decay as a function noncommutative scale A=1/4/|8| where 6=(623,031,12)."," In \cite{moj}, we have calculated the dependence of the charged lepton spin analyzing power in the top quark decay as a function noncommutative scale $\Lambda = 1/\sqrt{|\vec{\theta}|}$ where $\vec{\theta}=(\theta_{23},\theta_{31},\theta_{12})$."
+ Fig., Fig.
+ 3. presents the charged lepton spin analyzing power versus noncommutativity scale., \ref{alphanc} presents the charged lepton spin analyzing power versus noncommutativity scale.
+" Within the allowed region of noncommutativity scale (A21 TeV) [0],, the deviation from the SM prediction is less than 0.01%."," Within the allowed region of noncommutativity scale $\Lambda \gtrsim 1$ TeV) \cite{moj2}, the deviation from the SM prediction is less than $0.01\%$."
+" Finally, Table 1 shows the deviations which the charged lepton and the b-quark spin analyzing powers can receive from different sources."," Finally, Table \ref{results} shows the deviations which the charged lepton and the $b$ -quark spin analyzing powers can receive from different sources."
+ As it can be seen within the SM a variation of 2 GeV/c? of the top quark mass leads to 2.3% change in the b-quark spin analyzing power and no change to αι., As it can be seen within the SM a variation of 2 $^2$ of the top quark mass leads to $2.3\%$ change in the $b$ -quark spin analyzing power and no change to $\alpha_{l}$ .
+ The maximum deviation with the size of 5% occurs due to the NLO QCD correction for the bottom quark., The maximum deviation with the size of $5\%$ occurs due to the NLO QCD correction for the bottom quark.
+"model,",model.
+ The bottom panel of Figure 5 shows an equivalent plot for the case where the contributed bv the water emission is smaller. JJs. the strength of the flux attributed (ο component (4) in the Eisner (2007) model.," The bottom panel of Figure 5 shows an equivalent plot for the case where the pseudo-continuum contributed by the water emission is smaller, Jy, the strength of the flux attributed to component (4) in the Eisner (2007) model."
+ In this case as well. the MWC 480 spectrum shows less structure (han in the model.," In this case as well, the MWC 480 spectrum shows less structure than in the model."
+ Thus. water emission seems unlikely be present at the level needed to account for the flux attributed to either the hot compact component or the water component of the Eisner (2007) model.," Thus, water emission seems unlikely be present at the level needed to account for the flux attributed to either the hot compact component or the water component of the Eisner (2007) model."
+ A possible caveat is that the water emission reported by Eisner (2007) may be time variable., A possible caveat is that the water emission reported by Eisner (2007) may be time variable.
+ The A-band flux of MWC 480 is known to vary with time by a modest. amount(159: Sitko οἱ 22008: de Winter et 22001)., The $K$ -band flux of MWC 480 is known to vary with time by a modest amount; Sitko et 2008; de Winter et 2001).
+ Llowever. the strength of the A-band continuum in our spectrum is similar to that reported by Eisner (2007). which does not suggest a large spectral variation between (he two epochs of observation.," However, the strength of the $K$ -band continuum in our spectrum is similar to that reported by Eisner (2007), which does not suggest a large spectral variation between the two epochs of observation."
+ Another possible caveat is that V1331 (νο is not a good template for the water emission in MWC! 480., Another possible caveat is that V1331 Cyg is not a good template for the water emission in MWC 480.
+ The water emission in \IWC 480 may be much more highly optically thick and line blanketed. involving much higher water column densities than (hose considered by Eisner (2007) or than is present in the V1331 Cvg spectrum.," The water emission in MWC 480 may be much more highly optically thick and line blanketed, involving much higher water column densities than those considered by Eisner (2007) or than is present in the V1331 Cyg spectrum."
+ The model of water emission used by Eisner (2007) to fit the interferometer observations of 4480 is far [rom being line blanketed., The model of water emission used by Eisner (2007) to fit the interferometer observations of 480 is far from being line blanketed.
+ It is in fact more optically thin than the water enission observed in (he V1331 (νο spectrum., It is in fact more optically thin than the water emission observed in the V1331 Cyg spectrum.
+ The model parameters used by Eisner (2007). a temperature of KIX and a column density of 1.2xLOMem.7. do a reasonable job reproducing the relative [Iuxes of manv of the water lines in the region immediately blueward of the CO overtone bandhbeacd. with the important exception that the lines at 2.2017 and 2.2928jm are factors of 2:3 too strong compared to the weaker water lines.," The model parameters used by Eisner (2007), a temperature of K and a column density of $1.2\times 10^{19}\persqcm$, do a reasonable job reproducing the relative fluxes of many of the water lines in the region immediately blueward of the CO overtone bandhead, with the important exception that the lines at 2.2917 and $2.2928\micron$ are factors of 2–3 too strong compared to the weaker water lines."
+ This is because the water emission in the Eisner (2007) model is optically thin. whereas (he relative fluxes of the water lines in the V1331 Cve spectrum indicate more optically thick emission.," This is because the water emission in the Eisner (2007) model is optically thin, whereas the relative fluxes of the water lines in the V1331 Cyg spectrum indicate more optically thick emission."
+ A more detailed study of possible water emission in MWC 480 would benefit [rom the development of a reliable water line list for the short wavelength end of the Ax-band., A more detailed study of possible water emission in MWC 480 would benefit from the development of a reliable water line list for the short wavelength end of the $K$ -band.
+ This is a subject for future work., This is a subject for future work.
+ V1331 Cve shows a rich spectrum of water emission in the A-band., V1331 Cyg shows a rich spectrum of water emission in the $K$ -band.
+ In comparison. the {ας spectrum of MWC 480 shows little emission from water or anv other spectral Ines.," In comparison, the $K$ -band spectrum of MWC 480 shows little emission from water or any other spectral lines."
+ The non-delection of water emission is consistent wilh the absence of CO overtone emission from MWC 480: all sources of A-band water emission reported in the literature also show, The non-detection of water emission is consistent with the absence of CO overtone emission from MWC 480: all sources of $K$ -band water emission reported in the literature also show
+IRC+10216 is the prototypical late-stage. carbon-rich asvimptotic giant. branch (AGB) star.,"IRC+10216 is the prototypical late-stage, carbon-rich asymptotic giant branch (AGB) star."
+ This star is extremely well-studied. due to its close proximity. and large mass loss rate: the star is losing mass al a rale of ~3x10.7?M. vr.| (Glassgold 1996). producing a dense. dustv. circinstellar envelope. well shielded [rom interstellar ultraviolet. (SUV) radiation.," This star is extremely well-studied, due to its close proximity and large mass loss rate; the star is losing mass at a rate of $\sim 3 \times 10^{-5}\,$ $_{\odot}\,{\rm yr^{-1}}$ (Glassgold 1996), producing a dense, dusty circumstellar envelope, well shielded from interstellar ultraviolet (ISUV) radiation."
+ The dense. shielded environment of IRC+10216's envelope is home to a rich. cireumstellar chemistry.," The dense, shielded environment of IRC+10216's envelope is home to a rich circumstellar chemistry."
+ To date. more than 50 molecules have been detected around IRC+10216.," To date, more than 50 molecules have been detected around IRC+10216."
+ Due to the extreme carbon-rich nature of the source (C/O =1.4) (Glassgold 1996). the detection ol any oxvgen-bearing molecules other than CO and small amounts of SiO and was entirely unexpected.," Due to the extreme carbon-rich nature of the source (C/O $\gtrsim 1.4$ ) (Glassgold 1996), the detection of any oxygen-bearing molecules other than CO and small amounts of SiO and $^+$ was entirely unexpected."
+ Recently. however. Melnick οἱ al. (," Recently, however, Melnick et al. ("
+2001) reported the detection of an emission feature al 556.936 GlIz. attributed to the lyy—lop transition of water vapor.,"2001) reported the detection of an emission feature at $556.936\,$ GHz, attributed to the $1_{10}-1_{01}$ transition of water vapor."
+ The detection of water vapor in such a carbon-rich. ciremustellar envelope was interpreted by Melnick et al. (, The detection of water vapor in such a carbon-rich circumstellar envelope was interpreted by Melnick et al. (
+2001) as evidence for the existence of an extrasolar cometary svstenm. analogous to the Solar Systems Ixuiper Bell. in orbit around IRO+10216.,"2001) as evidence for the existence of an extrasolar cometary system, analogous to the Solar System's Kuiper Belt, in orbit around IRC+10216."
+ In this svsten. (he Iuminositv of the central star has increased dramatically due to the later stages of sequence stellar evolution. causing the iev bodies in the Ixuiper Bell analog to vaporize (Ford Neuleld 2001) and produce the water vapor observed by Melnick et al. (," In this system, the luminosity of the central star has increased dramatically due to the later stages of post-main sequence stellar evolution, causing the icy bodies in the Kuiper Belt analog to vaporize (Ford Neufeld 2001) and produce the water vapor observed by Melnick et al. ("
+2001).,2001).
+ OI possible concern for the cometary hypothesis is the existence of a peculiar subclass of carbon stars which exhibit OIL and II4O maser emission. as well as silicate emission features (e.g. Engels 1994).," Of possible concern for the cometary hypothesis is the existence of a peculiar subclass of carbon stars which exhibit OH and $_2$ O maser emission, as well as silicate emission features (e.g. Engels 1994)."
+ These features would normally be unexpected in carbon stars. for the reasons explained above.," These features would normally be unexpected in carbon stars, for the reasons explained above."
+ This subclass also shows a number of other features which distinguish peculiar carbon stars from normal carbon stars. including relatively high O/C ralios. IRAS colors compatible with low mass loss rate oxvgen stars. and depleted s-process element abundances relative to other carbon stars.," This subclass also shows a number of other features which distinguish peculiar carbon stars from normal carbon stars, including relatively high $^{13}$ $^{12}$ C ratios, IRAS colors compatible with low mass loss rate oxygen stars, and depleted s-process element abundances relative to other carbon stars."
+ These peculiar carbon stus may represent objects in the transition stage between oxvgen star and carbon star. and (heir photospheric chemistry is likely to be out. of thermochenical equilibrium. or they may. be members of binary pairs of one oxveen star and one carbon star (Wallerstein Knapp 1998).," These peculiar carbon stars may represent objects in the transition stage between oxygen star and carbon star, and their photospheric chemistry is likely to be out of thermochemical equilibrium, or they may be members of binary pairs of one oxygen star and one carbon star (Wallerstein Knapp 1998)."
+ In any case. IRC+10216 clisplavs none of the ancillary characteristics associated wilh these peculiar carbon stars. aside from. OIL and 115O emission.," In any case, IRC+10216 displays none of the ancillary characteristics associated with these peculiar carbon stars, aside from OH and $_2$ O emission."
+" Due to the large C/O ratio in IRC+10216 and the [aet that it is not a member of a binary svstem. as well as the lack of ancillary characteristics. we exclude the possibility (hat OLL or Ε.Ο emission trom IRC+10216 arises from membership in this peculiar class of carbon stars,"," Due to the large C/O ratio in IRC+10216 and the fact that it is not a member of a binary system, as well as the lack of ancillary characteristics, we exclude the possibility that OH or $_2$ O emission from IRC+10216 arises from membership in this peculiar class of carbon stars."
+ mince (he Melnick et al. (, Since the Melnick et al. (
+2001) result was based on the detection of a single line of water vapor. we became interested in finding a method of verifving the identification of water vapor with the emission feature αἱ 556.936 GlIz.,"2001) result was based on the detection of a single line of water vapor, we became interested in finding a method of verifying the identification of water vapor with the emission feature at $556.936\,$ GHz."
+ To this end. we realized that the water vapor emitted by iev bodies in the cometary svstem would be carried by the cireimstellar outflow," To this end, we realized that the water vapor emitted by icy bodies in the cometary system would be carried by the circumstellar outflow"
+"Whether the star is bound to the Galaxy highly depends on the Galactic potential adopted, in on the mass of the dark matter halo, as pointed out by particularAbadietal.(2009).","Whether the star is bound to the Galaxy highly depends on the Galactic potential adopted, in particular on the mass of the dark matter halo, as pointed out by \citet{2009ApJ...691L..63A}."
+. Our Mbalo is similar to values found in several recent studies., Our $M_{\rm halo}$ is similar to values found in several recent studies.
+" Wilkinson&Evans(1999) used 27 satellite galaxies and globular clusters, by assuming that they are bound and derived a total Galactic halo mass of Mya,~1.9*26x101”msun."," \citet{1999MNRAS.310..645W} used 27 satellite galaxies and globular clusters, by assuming that they are bound and derived a total Galactic halo mass of $M_{\rm halo}\sim1.9_{-1.7}^{+3.6}\times10^{12}$."
+. This value matches our derivation but has larger uncertainty., This value matches our derivation but has a larger uncertainty.
+" Sakamotoetal.(2003) used 11 satellite galaxies,a 137 globular clusters and 413 field horizontal branch stars to derive a total Galactic mass."," \citet{2003A&A...397..899S} used 11 satellite galaxies, 137 globular clusters and 413 field horizontal branch stars to derive a total Galactic mass."
+" The exclusion of from their sample would lower the total Galactic mass from 2.5705x10! tto a value of Mio~1.805x10?msun,, both again in excellent agreement with our finding."," The exclusion of from their sample would lower the total Galactic mass from $M_{\rm total}\sim2.5_{-1.0}^{+0.5}\times10^{12}$ to a value of $M_{\rm total}\sim1.8_{-0.7}^{+0.4}\times10^{12}$, both again in excellent agreement with our finding."
+" On the other hand, less than of the trajectories resulting from our MC simulations would be bound for the most likely mass of Xue et al.,"," On the other hand, less than of the trajectories resulting from our MC simulations would be bound for the most likely mass of Xue et al.,"
+" Mya,=1.0x10? ((grey shaded area in Fig. 4))"," $M_{\rm
+halo}=1.0\times10^{12}$ (grey shaded area in Fig. \ref{fig:veldistrib_J1539}) )."
+ Their RV study of ~2400 blue horizontal-branch stars — which includes J1539+0239 but lacks any proper motion measurements — therefore likely underestimate the Galactic halo mass., Their RV study of $\sim$ 2400 blue horizontal-branch stars – which includes J1539+0239 but lacks any proper motion measurements – therefore likely underestimate the Galactic halo mass.
+ We reported the quantitative spectral analysis of a high-velocity star from the sample of faint blue halo stars of Xueetal.(2008)., We reported the quantitative spectral analysis of a high-velocity star from the sample of faint blue halo stars of \citet{2008ApJ...684.1143X}.
+ J1539+0239 was confirmed to be a Population II blue horizontal branch star with a low metallicity of [Fe/H] =--2.0 and the characteristic enhancement of a-elements., J1539+0239 was confirmed to be a Population II blue horizontal branch star with a low metallicity of $[$ $]=-2.0$ and the characteristic enhancement of $\alpha$ -elements.
+" Hereby, we performed a NLTE analysis of a halo BHB star for the first time."," Hereby, we performed a NLTE analysis of a halo BHB star for the first time."
+" While the majority of the weak lines were confirmed to be formed close to LTE conditions, many of the stronger metal lines — which are important diagnostics at the spectral resolution achieved within theSDSS — show pronounced NLTE strengthening, with the differences between the derived LTE and NLTE abundances amounting to ddex to ddex typically."," While the majority of the weak lines were confirmed to be formed close to LTE conditions, many of the stronger metal lines – which are important diagnostics at the spectral resolution achieved within the – show pronounced NLTE strengthening, with the differences between the derived LTE and NLTE abundances amounting to dex to dex typically."
+" In addition to information on the chemical composition, the radial velocity, proper motion and spectroscopic distance were derived and a detailed kinematical analysis was performed."," In addition to information on the chemical composition, the radial velocity, proper motion and spectroscopic distance were derived and a detailed kinematical analysis was performed."
+" Carrying out kinematical numerical experiments using the Galactic of Allen&Santillan(1991) in order to obtain anpotential orbit of J1539+0239 gravitationally bound to the Milky Way, we found that the mass of the dark halo has to be at least Mpaio~1.725x1013 (absolute uncertainties from extrema in MC error propagation)."," Carrying out kinematical numerical experiments using the Galactic potential of \citet{1991RMxAA..22..255A} in order to obtain an orbit of J1539+0239 gravitationally bound to the Milky Way, we found that the mass of the dark halo has to be at least $M_{\rm halo}\sim1.7_{-1.1}^{+2.3}\times10^{12}$ (absolute uncertainties from extrema in MC error propagation)."
+ This masslimit is in good agreement with several previous studies (Wilkinson&Evans1999;Sakamotoetal.2003;Abadi 2009)..," This masslimit is in good agreement with several previous studies \citep{1999MNRAS.310..645W,2003A&A...397..899S,2009ApJ...691L..63A}."
+" However, the significantly lower most likely mass value of Xueet(2008) is consistent with our analysis only at a level of4%,, i.e. it underestimates the Galactic dark halo mass."," However, the significantly lower most likely mass value of \cite{2008ApJ...684.1143X} is consistent with our analysis only at a level of, i.e. it likely underestimates the Galactic dark halo mass."
+" We likelyconclude, that if the kinematics of a halo star is extraordinary enough, and the errors within the analysis are small, even star alone can provide a significant lower limit to the dark matter halo mass, and to the total mass of the Milky Way ddisk), here Miota=x10?msun.."," We conclude, that if the kinematics of a halo star is extraordinary enough, and the errors within the analysis are small, even star alone can provide a significant lower limit to the dark matter halo mass, and to the total mass of the Milky Way disk), here $M_{\rm total}\ge1.8_{-1.1}^{+2.3}\times10^{12}$."
+" The determining factor is that full1.8775 kinematic information is available, as it will become routine in the era of the Gaia space mission, at much higher precision."," The determining factor is that full kinematic information is available, as it will become routine in the era of the Gaia space mission, at much higher precision."
+ A.T. acknowledges funding by the Deutsche Forschungsgemeinschaft (DFG) through grant HE1356/45-1., A.T. acknowledges funding by the Deutsche Forschungsgemeinschaft (DFG) through grant HE1356/45-1.
+" Travel to the DSAZ (Calar Alto, Spain) was supported by the DFG under grant HE1356/50-1."," Travel to the DSAZ (Calar Alto, Spain) was supported by the DFG under grant HE1356/50-1."
+ We are very grateful to Stephan Geier for stimulating discussions and advice., We are very grateful to Stephan Geier for stimulating discussions and advice.
+ Our thanks go to S. Miilller and T. Kupfer for observing and reducing the data from DSAZ., Our thanks go to S. Mülller and T. Kupfer for observing and reducing the data from DSAZ.
+" Funding for the and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England."," Funding for the and -II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England."
+" The Web Site is Sloan,, (TWIN).."," The Web Site is , ."
+As meutioucl in the Introduction. the quantity to be measured is tlje radius of the surface at which the local temperature equals the solar effective temperature.,"As mentioned in the Introduction, the quantity to be measured is the radius of the surface at which the local temperature equals the solar effective temperature."
+ Evideutlv. this radius is related to Dy in Eq. (," Evidently, this radius is related to $D_0$ in Eq. ("
+2). but the relationship 1s not a simple one: it depeuds upou the radiatiou traister iu the outer solar atinospliere. and upon the hejiwvior of the FETD lib definition.,"2), but the relationship is not a simple one; it depends upon the radiation transfer in the outer solar atmosphere, and upon the behavior of the FFTD limb definition."
+ To inter the coxrect radius from the observations. one nist use a oivsicallv-based. inodel of the solar atinosphere to cdeulate the eierseut mteusitv as a function of distatice from the center of the solar disk. and then conite the location on this brightness profile that woud be identified as the edge by the ΕΕΤΙ).," To infer the correct radius from the observations, one must use a physically-based model of the solar atmosphere to calculate the emergent intensity as a function of distance from the center of the solar disk, and then compute the location on this brightness profile that would be identified as the edge by the FFTD."
+ We calculated the limb intensityby iutegratiug the equation of trausfer along ravs through an assumed spherically sviuiuetrical solar atmosphere., We calculated the limb intensity by integrating the equation of transfer along rays through an assumed spherically symmetrical solar atmosphere.
+ Since the observations were carried out in a relatively narrow wavelength region around 500 1uu. we considered simply the monochromatic coutimun intensity at this waveleneth.," Since the observations were carried out in a relatively narrow wavelength region around 800 nm, we considered simply the monochromatic continuum intensity at this wavelength."
+ We have used two models of the solar atiuospliere., We have used two models of the solar atmosphere.
+ Model l was kindly computed w Roduev Medupe with the ATLASS code Usurucz 1993)., Model 1 was kindly computed by Rodney Medupe with the ATLAS9 code (Kurucz 1993).
+ Model 2 was obtained as an average of a ivdrodyvinauaical simulation of convection in the upper part of he solar convection gone and lower atmosphere. as described by. eg. Stein Nordlund (1989) ut wili pdated plysics (Trampedach 1997): the average was performed at constant monochromatic optical depth. at S00 niu.," Model 2 was obtained as an average of a hydrodynamical simulation of convection in the upper part of the solar convection zone and lower atmosphere, as described by, e.g., Stein Nordlund (1989) but with updated physics (Trampedach 1997); the average was performed at constant monochromatic optical depth, at 800 nm."
+ The opacity was computed frou he ATLAS data in both cases., The opacity was computed from the ATLAS data in both cases.
+ For Model 1 the source function Sy was obtained from the ATLAS code. and hence allowed for mild departure roni LTE.," For Model 1 the source function $S_\lambda$ was obtained from the ATLAS code, and hence allowed for mild departure from LTE."
+ For Model 2 LTE was assumed. so that Sy=By. the Planck function.," For Model 2 LTE was assumed, so that $S_\lambda = B_\lambda$, the Planck function."
+ To simulate effects of seciug. we convolved the intensity with a gaussia. with full width at halt maxim (FWIHALD) specified in ar€ sec and converted to linear cüstauce at J AU.," To simulate effects of seeing, we convolved the intensity with a gaussian, with full width at half maximum (FWHM) specified in arc sec and converted to linear distance at 1 AU."
+" After tlis convolution. we integrated the intensities over pixeIs corresponding to 1"" at 1 AU. to match the observed. iuteusities;"," After this convolution, we integrated the intensities over pixels corresponding to $1 \arcsec$ at 1 AU, to match the observed intensities."
+" We folded the pixel-weighted intensities with the FFTD weights over the five differeu windows described iu section 2. and carne out the subsequent analysis to determine the linib position. thro,eh extrapolation to zero Window width. as for the observations."," We folded the pixel-weighted intensities with the FFTD weights over the five different windows described in section 2, and carried out the subsequent analysis to determine the limb position, through extrapolation to zero window width, as for the observations."
+" The results of applying the observational procedure to the computec pixol-averaged menusities for Models laud 2. assuniue 6"" sccing are shown in Fig."," The results of applying the observational procedure to the computed pixel-averaged intensities for Models 1 and 2, assuming $6 \arcsec$ seeing are shown in Fig."
+ 1., 1.
+ The observed variation of diameter w3th window width aud that calculated from the modes agree well. especially or Model 1.," The observed variation of diameter with window width and that calculated from the models agree well, especially for Model 1."
+ Of particular interest :we the extrapolations ο zero Window width. corresponding to the observed inb position measured relative to he nominal photosphere ofthe models: weobtain 0.17950.Maa aud 0.51631αι or Models 1 aud 2.," Of particular interest are the extrapolations to zero window width, corresponding to the observed limb position measured relative to the nominal photosphere of the models; we obtain $0.47950 \; \Mm$ and $0.51634 \; \Mm$ for Models 1 and 2."
+ We have tested the sensitivity* of the results to various assuniptious in the calculation., We have tested the sensitivity of the results to various assumptions in the calculation.
+ Replacing the rue source function Sy bv By for Model 1 changes he lim position x onmuch less than 0.001 Maa: thus he asstunption of LTE in Model 2 is not a significant source of error., Replacing the true source function $S_\lambda$ by $B_\lambda$ for Model 1 changes the limb position by much less than $0.001$ Mm; thus the assumption of LTE in Model 2 is not a significant source of error.
+ Using au assumed. secing of less than chanees the limb position by less than 0.01Maa. confirming the mseusitivitv of tlie FFTD to effects of secius.," Using an assumed seeing of less than changes the limb position by less than $0.01 \Mm$, confirming the insensitivity of the FFTD to effects of seeing."
+ However. the difference between the two models obviously reais a source! of so1ue concern.," However, the difference between the two models obviously remains a source of some concern."
+The description of close binary svstenus is usually based ou the Roche model which defines the shape of a binary component distorted by tidal and rotational forces.,The description of close binary systems is usually based on the Roche model which defines the shape of a binary component distorted by tidal and rotational forces.
+ Iu re framework of the Roche model one assmues tha je binary conmiponeuts (the primary and the secondary) 4ther are point masses. or are corotating and have a PAyherically syauunnetrie inass distribttion Dresoective of ici proxinütv or mass ratio (Noped 1959. 16NN)," In the framework of the Roche model one assumes that the binary components (the primary and the secondary) either are point masses, or are corotating and have a spherically symmetric mass distribution irrespective of their proximity or mass ratio (Kopal 1959, 1978)."
+ Iu a send¢etached system. one of the componcus fills its critica equipoteutial lobe defined w the potential of ie inner Lagrangkuir point. and which determines the asin extent of a star ina close binary.," In a semi–detached system, one of the components fills its critical equipotential lobe defined by the potential of the inner Lagrangian point, and which determines the maximum extent of a star in a close binary."
+ This is the so Roche lobe within the Roche model., This is the so--called Roche lobe within the Roche model.
+ Cataclvsimic variabes (hereafter CVs). composed of a white dwarf as the primary and a lowmass star or a brown dwarf as the secondary. belong to this type of svstems: the SOCOMary fills its criical lobe and transfers mass towards the primary.," Cataclysmic variables (hereafter CVs), composed of a white dwarf as the primary and a low–mass star or a brown dwarf as the secondary, belong to this type of systems: the secondary fills its critical lobe and transfers mass towards the primary."
+ When applying the Roche moel to wobblems of binary evolution OMe ludakes implicitly the ollowius assuniptions (among others)., When applying the Roche model to problems of binary evolution one makes implicitly the following assumptions (among others).
+ First. the Roche poteutial is a good approxiuatio Loft the rue potential that one would obtain by solving the Poisson equatio1.," First, the Roche potential is a good approximation of the true potential that one would obtain by solving the Poisson equation."
+ Second. the effects of tidal aud rotational forces on the internal structure of the star are negligible. ic. tthat they result iu only. siuall correctiois conipared to stellar modes assiunainue spherical eeonetrv.," Second, the effects of tidal and rotational forces on the internal structure of the star are negligible, i.e. that they result in only small corrections compared to stellar models assuming spherical geometry."
+ Third. that for he purpose of evolutionary ConttaΊος. involving onedimensional stellar τοσο] ιο loICfilling star iiw be replaced by a spherical star ο fthe same volume.," Third, that for the purpose of evolutionary computations involving one–dimensional stellar models the lobe–filling star may be replaced by a spherical star of the same volume."
+ This is antamount to assumniue tha 1011 itidal and rotational forces change the shape of a ar they* leave its volue invariant., This is tantamount to assuming that though tidal and rotational forces change the shape of a star they leave its volume invariant.
+ The radius of the (yefilliis star then only depends on the geometry of je system. aud can be calculated by means of simple nalvtical fits (Paczyvisski 1971. Egeletou 1983).," The radius of the lobe–filling star then only depends on the geometry of the system, and can be calculated by means of simple analytical fits (Paczyńsski 1971, Eggleton 1983)."
+ The main )urpose of the present paper is to examine in some detai je third aud to some extent also the secoud of the above assuniptious. both of which have so far not been tested.," The main purpose of the present paper is to examine in some detail the third and to some extent also the second of the above assumptions, both of which have so far not been tested."
+ Recent 3D simulations (Rezzolla et. al., Recent 3D simulations (Rezzolla et al.
+ 2001: Motl Frank. priv.," 2001; Motl Frank, priv."
+ couunu.), comm.)
+ confirm that. at least im the case of a scuni-detached svsteui the Roche potential is a good approximation if the olc-filliug star is sutiicicutly ceutrally coudensed. ie. if the cfective polytropic index is N> 3/2.," confirm that, at least in the case of a semi-detached system, the Roche potential is a good approximation if the lobe-filling star is sufficiently centrally condensed, i.e. if the effective polytropic index is $N \simgr$ 3/2."
+ The analvsis of RezzMa et al. (, The analysis of Rezzolla et al. (
+2001) is based on numerical models of semi-cdeached binaries that account for the finite size of the secudaryv star. thus rolaxiis the first assumption iuherenu in the Roche uocdel.,"2001) is based on numerical models of semi-detached binaries that account for the finite size of the secondary star, thus relaxing the first assumption inherent in the Roche model."
+ With the validity of this approximation for the dcTerninatkn Q: the potential. they also show that such effects haxlv affect eravitationa quadiupole racdiatlon.," With the validity of this approximation for the determination of the potential, they also show that such effects hardly affect gravitational quadrupole radiation."
+ Aloreover. a οςonuparison betweuu the augular nonueutimw loss axd nüass-transfer timescales predicted by he Roche iuccl xd their uunerica models shows smal differences.," Moreover, a comparison between the angular momentum loss and mass-transfer timescales predicted by the Roche model and their numerical models shows small differences."
+" They ius conclude that finite size effects camnot account for tic 11jisuatceh between the observed munimi period P4, at ) uiu of CV svsclus and the theoretica value fau.", They thus conclude that finite size effects cannot account for the mismatch between the observed minimum period $\pmin$ at 80 min of CV systems and the theoretical value $\pturn$ .
+ The later ds indeed ~ shorter than the observe value. ac€odi1ο to recent caleulatious base On mnuiprovec scllar plivsics (see νους Doaraffe 190€)).," The latter is indeed $\sim$ shorter than the observed value, according to recent calculations based on improved stellar physics (see Kolb Baraffe 1999)."
+ Since Rezzolla e al. (, Since Rezzolla et al. (
+2) VL)«o not cousider thermal reaxalon effects in tjor caclations. they cau only determine a ciffereutia CCxrectk1 ifo Ius2 when eoiug from Roche nodel to selt-CCπρ1 pochtial.,"2001) do not consider thermal relaxation effects in their calculations, they can only determine a differential correction to $\pturn$ when going from Roche model to self-consistent potential."
+ ILowever. iu cdoine x» the secoix aud third assniptious iieutioned above remain uutested.," However, in doing so, the second and third assumptions mentioned above remain untested."
+ The τιlai purpose of our paper is thus to explore the CCMiscποices of making these Wo asstuuptions., The main purpose of our paper is thus to explore the consequences of making these two assumptions.
+" Qur nain coal is to deternüne quantitatively the departure from spherical siunetry of the secondary du Xquietached binaries and to analyse the COMSCQUCLCCS onu the amass trausfer rates and orbital period in CV απο»,", Our main goal is to determine quantitatively the departure from spherical symmetry of the secondary in semi-detached binaries and to analyse the consequences on the mass transfer rates and orbital period in CV systems.
+ We use siioothied paricle lvdvodvuamiucs (SPIT) techniques to study equilibrini coufiguratious of seni-detached binaries and estimate for different mass ratios the geometrical deformation of the secoudary as it fills its critical lobe., We use smoothed particle hydrodynamics (SPH) techniques to study equilibrium configurations of semi-detached binaries and estimate for different mass ratios the geometrical deformation of the secondary as it fills its critical lobe.
+ The nunueneal models and results are described in 82., The numerical models and results are described in 2.
+ Iu 83. we analyse some of the consequences of the tidal aud rotational forces on the secular evolution of the low mass donor on grounds of models constructed bvIvolh Daraffe (1999) aud Baratte νο (2000).," In 3, we analyse some of the consequences of the tidal and rotational forces on the secular evolution of the low mass donor on grounds of models constructed byKolb Baraffe (1999) and Baraffe Kolb (2000)."
+The GMRT (Swarup et al.,The GMRT (Swarup et al.
+" 1991) observations of HIZSSO03 (RA (2000): 07"" 0029,37. DEC(2000); 04712/30” ) were conducted on 23 Aug 2004."," 1991) observations of HIZSS003 (RA (2000): $^h$ $^m$ $^s$, DEC(2000): $-{04}^{\circ} 12^\prime 30^{\prime\prime}$ ) were conducted on 23 Aug 2004."
+ An observing bandwidth of | MHz centered at 1419.1 MHz (which corresponds to a heliocentric velocity of 290 5» was used., An observing bandwidth of 1 MHz centered at 1419.1 MHz (which corresponds to a heliocentric velocity of 290 ) was used.
+ The band was divided into [28 spectral channels. giving a channel spacing of 1.65!.," The band was divided into 128 spectral channels, giving a channel spacing of 1.65."
+. Flux calibration was done using scans on the standard calibrators 3C147 and 3C286. which were observed at the start and end of the observing run.," Flux calibration was done using scans on the standard calibrators 3C147 and 3C286, which were observed at the start and end of the observing run."
+ Phase calibration was done using 0744-064. which was observed once every 40 minutes.," Phase calibration was done using 0744-064, which was observed once every 40 minutes."
+ Bandpass calibration was done in the standard way using 3C286., Bandpass calibration was done in the standard way using 3C286.
+ The total on-source time was ~ 4+ hours., The total on-source time was $\sim$ 4 hours.
+ The data was reduced using standard tasks in classic AIPS., The data was reduced using standard tasks in classic AIPS.
+" The GMRT has a hybrid configuration which simultaneously provides both high angular resolution (~2"" if one uses baselines between the arm antennas) as well as sensitivity to extended emission (from baselines between the antennas in the central array).", The GMRT has a hybrid configuration which simultaneously provides both high angular resolution $\sim 2^{''}$ if one uses baselines between the arm antennas) as well as sensitivity to extended emission (from baselines between the antennas in the central array).
+ Data cubes were therefore made at various resolutions including 490 39.28.ο 26.23.IN LIS 11.8.6. and 4053 .using uniform weighting.," Data cubes were therefore made at various resolutions including $^{''}\times 39^{''}$, $^{''}\times 26^{''}$, $^{''}\times 18^{''}$, $^{''}\times 11^{''}$, $^{''}\times 6^{''}$ and $4^{''}\times3^{''}$, using uniform weighting."
+ RMS noise per channel for these resolution is 2.2 mJy. 2.0 mJy. 1.8 mJy. 1.6 mJy. 1.4 mJy and 1.2 mJy respectively.," RMS noise per channel for these resolution is 2.2 mJy, 2.0 mJy, 1.8 mJy, 1.6 mJy, 1.4 mJy and 1.2 mJy respectively."
+ All the data cubes. except 8).6 and4.37 were deconvolved using the task IMAGR.," All the data cubes, except $^{''}\times 6^{''}$ and $4^{''}\times3^{''}$, were deconvolved using the task IMAGR."
+ For the two highest resolution data cubes. the signal to noise ratio was too low for CLEAN to work reliably.," For the two highest resolution data cubes, the signal to noise ratio was too low for CLEAN to work reliably."
+ A continuum image was made using the average of the line free channels., A continuum image was made using the average of the line free channels.
+ No continuum was detected from the galaxy to a 30 flux limit of 1.0 mJy/beam (for a beam size of 2852026 , No continuum was detected from the galaxy to a $3\sigma$ flux limit of 1.0 mJy/beam (for a beam size of $28^{''}\times26^{''}$ ).
+4A high resolution continuum map GECX resolution) was also made to search for any compact continuum sources., A high resolution continuum map $4^{''}\times3^{''}$ resolution) was also made to search for any compact continuum sources.
+ The only continuum source of note detected is NVSS I070023-041255., The only continuum source of note detected is NVSS J070023-041255.
+" The HI column density (as derived from the 42.39 resolution image) along the line of sight to this source is 107"" atoms 7.", The HI column density (as derived from the $42^{''}\times 39^{''}$ resolution image) along the line of sight to this source is $\times~10^{20}$ atoms $^{-2}$.
+ A search for HI absorption. in the direction of this source. gave negative results at all resolutions.," A search for HI absorption, in the direction of this source, gave negative results at all resolutions."
+ The implied lower limit on the spin temperature of the gas (assuming a velocity width of 10 5j) is 723 K. For reference we note that NVSS J070023- lies towards the source that we call HIZSSOO3B below. and that it is close to. but not coincident with the HII region detected by Massey et al. (," The implied lower limit on the spin temperature of the gas (assuming a velocity width of 10 ) is 723 K. For reference we note that NVSS J070023-041255 lies towards the source that we call HIZSS003B below, and that it is close to, but not coincident with the HII region detected by Massey et al. ("
+2003).,2003).
+ It appears likely that it is a background source. with no connection to the HI emission.," It appears likely that it is a background source, with no connection to the HI emission."
+ The continuum source was subtracted using the task UVSUB., The continuum source was subtracted using the task UVSUB.
+ Channel maps of the HI emission at a resolution of 4os are shown in Fig. l.., Channel maps of the HI emission at a resolution of $42^{''}\times39^{''}$ are shown in Fig. \ref{fig:cube}.
+ HI emission is spread over 63 channels and consists of two distinct sources. one spanning 39 channels and the other spanning 26 channels.," HI emission is spread over 63 channels and consists of two distinct sources, one spanning 59 channels and the other spanning 26 channels."
+ At this spatial resolution some channels show HI emission connecting the two sources. however it is not clear whether this is due to beam smearing.," At this spatial resolution some channels show HI emission connecting the two sources, however it is not clear whether this is due to beam smearing."
+ A HI feature connecting the two sources is also seen in ονο.26 and 23.18” resolution data cubes., A HI feature connecting the two sources is also seen in $^{''}\times 26^{''}$ and $^{''}\times 18^{''}$ resolution data cubes.
+ However. at higher resolutions no such connecting emission is seen in the channel maps.," However, at higher resolutions no such connecting emission is seen in the channel maps."
+ Further. as discussed in more detail below. the velocity field of the bigger source does not appear to be particularly disturbed. and neither source shows signs of two armed tidal distortions.," Further, as discussed in more detail below, the velocity field of the bigger source does not appear to be particularly disturbed, and neither source shows signs of two armed tidal distortions."
+ It is possible that the connecting emission seen in Fig., It is possible that the connecting emission seen in Fig.
+ |. is due to beam smearing., \ref{fig:cube} is due to beam smearing.
+ In order to disentangle the HI emission one spectral cube was made for each galaxy in which emission from the other galaxy was blanked out., In order to disentangle the HI emission one spectral cube was made for each galaxy in which emission from the other galaxy was blanked out.
+ In the case of channel maps which showed connecting HI. the blanking was done midway between the two sources.," In the case of channel maps which showed connecting HI, the blanking was done midway between the two sources."
+ The Fig. 2[[A]&[B], The Fig. \ref{fig:mom0}[
+ resolution made from these blanked cubes., resolution made from these blanked cubes.
+" In the rest of the paper. we refer to the the bigger (eastern) galaxy as HIZSSOO3A and the smaller Qvestern) one as HIZSSOO3B. The entire HI distribution will be referred as the ""HIZSSO03 system”."," In the rest of the paper, we refer to the the bigger (eastern) galaxy as HIZSS003A and the smaller (western) one as HIZSS003B. The entire HI distribution will be referred as the “HIZSS003 system""."
+ The sources HIZSSO03A and HIZSSOO3B correspond to the main HI peak and the secondary unresolved peak in the VLA map of Massey et al. (, The sources HIZSS003A and HIZSS003B correspond to the main HI peak and the secondary unresolved peak in the VLA map of Massey et al. (
+2003).,2003).
+ The combination of poor spatial (~607) and velocity (~10 1) resolution of the VLA observations prevented Massey et al. (, The combination of poor spatial $\sim 60^{\prime\prime}$ ) and velocity $\sim$ 10 ) resolution of the VLA observations prevented Massey et al. (
+2003) from separating the two galaxies. although the near-IR VLT images do show two separate stellar concentrations. i.e. one for each galaxy.,"2003) from separating the two galaxies, although the near-IR VLT images do show two separate stellar concentrations, i.e. one for each galaxy."
+ Fig. 2[[, Fig. \ref{fig:mom0}[ [
+"C] shows the high resolution HI map (8...6"" resolution) of HIZSSO03.",C] shows the high resolution HI map $''\times6''$ resolution) of HIZSS003.
+ The more diffuse emission is resolved out. and the remaining emission from the two galaxies can be disentangled without having to resort to channel by channel blanking.," The more diffuse emission is resolved out, and the remaining emission from the two galaxies can be disentangled without having to resort to channel by channel blanking."
+ As can be seen in the Fig. 2..," As can be seen in the Fig. \ref{fig:mom0},"
+ the HI distribution in both the galaxies is clumpy. with three main peaks seen in the HI distribution of HIZSSO03A. whereas the HI distribution of HIZSSOO3B is resolved into two peaks.," the HI distribution in both the galaxies is clumpy, with three main peaks seen in the HI distribution of HIZSS003A, whereas the HI distribution of HIZSS003B is resolved into two peaks."
+ No signature of tidal interaction is evident in the HI distribution of either galaxy., No signature of tidal interaction is evident in the HI distribution of either galaxy.
+ The HII region detected in the HIZSSO03 system is located. close to one of the peaks of the HI distribution in HIZSSOO03B (shown by a cross in Fig. 2[[, The HII region detected in the HIZSS003 system is located close to one of the peaks of the HI distribution in HIZSS003B (shown by a cross in Fig. \ref{fig:mom0}[ [
+C]j.,C]).
+ The Ha emission is approximately aligned with the HI contours of the galaxy (i.e. from northwest to southeast). and its heliocentric velocity (335315 ‘Massey et al.," The $\alpha$ emission is approximately aligned with the HI contours of the galaxy (i.e. from northwest to southeast), and its heliocentric velocity $\pm$ 15 – Massey et al."
+ 2003). matches within the error bars with the systemic velocity of 322.6 + 1.4 ‘for HIZSS003B derived from the HI global profile (see below).," 2003), matches within the error bars with the systemic velocity of 322.6 $\pm$ 1.4 for HIZSS003B derived from the HI global profile (see below)."
+ Fig., Fig.
+ 3 shows the global HI emission profiles of the two galaxies obtained from the 42”39 resolution data cubes., \ref{fig:spectra} shows the global HI emission profiles of the two galaxies obtained from the $42{''}\times39^{''}$ resolution data cubes.
+ As discussed above. emission from one galaxy was blanked before obtaining the HI protile for the other one.," As discussed above, emission from one galaxy was blanked before obtaining the HI profile for the other one."
+ Gaussian fits to the HI profiles give systemic velocities of 288.0+2.5kms and 322.6 + L4 !for HIZSSO03A and HIZSS003B respectively., Gaussian fits to the HI profiles give systemic velocities of $\pm$ 2.5 and 322.6 $\pm$ 1.4 for HIZSS003A and HIZSS003B respectively.
+" The corresponding velocity widths at of peak emission are 55 ""and28+. while the integrated fluxes are 20.0-Ε32.1 Jy 'and 3,8E 0.3Jys."," The corresponding velocity widths at of peak emission are 55 and 28, while the integrated fluxes are $20.9\pm2.1$ Jy and $3.8\pm0.3$ Jy."
+".. The HI masses corresponding to these integrated flues are L4 10M. and 2.610""AL..", The HI masses corresponding to these integrated flues are 1.4 $\times10^7 \rm{M_\odot}$ and 2.6$\times10^6 \rm{M_\odot}$.
+ The combined flux of both galaxies is 24.7 Jy which is in excellent agreement with the value of 24.9 Jy obtained from the VLAobservations by Massey et al. (," The combined flux of both galaxies is 24.7 Jy, which is in excellent agreement with the value of 24.9 Jy obtained from the VLAobservations by Massey et al. ("
+2003).,2003).
+ However both these values are somewhat lower than the flux integral of ~32 Jy estimated from the single dish observations by Henning et al. (," However both these values are somewhat lower than the flux integral of $\sim$ 32 Jy, estimated from the single dish observations by Henning et al. ("
+2000).,2000).
+" The HI diameter of the two valaxies. measured at a level of ~10°"" atoms 7 (from the 42"".30"" images) are ~6.5! (3.2 Kpor and ~3” (1.5 kpe)."," The HI diameter of the two galaxies, measured at a level of $\sim10^{19}$ atoms $^{-2}$ (from the ${''}\times39{''}$ images) are $\sim6.5^\prime$ (3.2 kpc) and $\sim3^\prime$ (1.5 kpc)."
+ Fig. 4[[A]&[B] 4[f[, Fig. \ref{fig:mom1}[ [
+A]) is regular and a large scale velocity gradient. consistent with systematic rotation. is seen across the galaxy.,"A]) is regular and a large scale velocity gradient, consistent with systematic rotation, is seen across the galaxy."
+ The velocity field is also mildly lopsided the isovelocity contours in the southern half of the galaxy are more curved than the northern half., The velocity field is also mildly lopsided $-$ the isovelocity contours in the southern half of the galaxy are more curved than the northern half.
+ Kinks are seen in the eastern isovelocity contours. close to the location of HIZSSOO3B. These kinks are more prominent in the higher resolution velocity fields (not shown).," Kinks are seen in the eastern isovelocity contours, close to the location of HIZSS003B. These kinks are more prominent in the higher resolution velocity fields (not shown)."
+ Rotation curves of HIZSSOO3A were derived using 39 26 .23 LS and 151l resolution velocity fields. using tilted ring fits.," Rotation curves of HIZSS003A were derived using $^{''}\times 39^{''}$ , $^{''}\times 26^{''}$ , $^{''}\times 18^{''}$ and $^{''}\times 11^{''}$ resolution velocity fields, using tilted ring fits."
+ The center and systemic velocity for the galaxy, The center and systemic velocity for the galaxy
+curve gains larger energy curing the shock front crossing.,curve gains larger energy during the shock front crossing.
+" Figure 4 shows the time history of the first adiabatic invariant of 2,/D7 for the same particles analyzed in Figure 3.", Figure 4 shows the time history of the first adiabatic invariant of $P_{\perp e}/B^2$ for the same particles analyzed in Figure 3.
+" (P,/B7) is averaged over (he upstream electrongvro-period of c,I. and is also normalized by the upstream value of P4/D2."," $\langle P_{\perp e}/B^2 \rangle$ is averaged over the upstream electrongyro-period of $\omega_{ce}^{-1}$, and is also normalized by the upstream value of $P_{\perp 0}/B_0^2$."
+ The dashed curve is almost constant. and we know that the electron motion is almost adiabatie curing the shock erossiug.," The dashed curve is almost constant, and we know that the electron motion is almost adiabatic during the shock crossing."
+ For the solid curve case (hat had (he strong interaction of electron with ESW. we find a sharp energy increase lor the time interval from (1—/o)oy=120 to 240. which is suggestive a acceleration.," For the solid curve case that had the strong interaction of electron with ESW, we find a sharp energy increase for the time interval from $(t-t_0)\omega_{pe} = 120$ to $240$, which is suggestive a non-adiabatic acceleration."
+ After (/οc240 we found the gradual energy increase in Figure 3b. but the baseline of the normalized invariant seems to be almost constant. except [or a sinusoidal oscillation.," After $(t-t_0)\omega_{pe} \sim 240$ we found the gradual energy increase in Figure 3b, but the baseline of the normalized invariant seems to be almost constant except for a sinusoidal oscillation."
+ This suggests that the electron heating is almost adiabatic after the electrons are convected. downstream. and the non-adiabatie acceleration occurs only when (he electron interacts with ESW in the shock transition region.," This suggests that the electron heating is almost adiabatic after the electrons are convected downstream, and the non-adiabatic acceleration occurs only when the electron interacts with ESW in the shock transition region."
+ Figure 5 shows the above two particle trajectories in (he c—y plane., Figure 5 shows the above two particle trajectories in the $x-y$ plane.
+ Both two trajectories are drawn from CX.Y)e(200οων.) ab ή=535 in the upstream region. and the electrons are convected towards positive ο along y=0.," Both two trajectories are drawn from $(X,Y) \sim (200~c/\omega_{pe},0)$ at $\omega_{pe}t = 535$ in the upstream region, and the electrons are convected towards positive $x$ along $y=0$."
+ The Larmor radii are very small in the upstream region because their thermal velocities are very small, The Larmor radii are very small in the upstream region because their thermal velocities are very small.
+ Note that this result was obtained by the one-dimensional. particle-in-cell simulation. but both positions o£.r and y ave caleulated.," Note that this result was obtained by the one-dimensional, particle-in-cell simulation, but both positions of $x$ and $y$ are calculated."
+ Also note that the 7 and y distances awe plotted on the dilferent scale in Figue 5., Also note that the $x$ and $y$ distances are plotted on the different scale in Figure 5.
+ Thev encounter the shock front region around (οών)=225., They encounter the shock front region around $X/(c/\omega_{pe}) = 225$.
+ The notation of the solid and dashed. curve is same as that in Figures 3 and +., The notation of the solid and dashed curve is same as that in Figures 3 and 4.
+ We find that the electron denoted by the solid curve are traveling along the y axis in the shock front region. it can be understood that the main energy gain of the electrons comes [rom the motional electric field.," We find that the electron denoted by the solid curve are traveling along the $y$ axis in the shock front region, it can be understood that the main energy gain of the electrons comes from the motional electric field."
+" In fact. the energv gain of electron from the motional electric field Ac,, ean be estimated as Y€E,,. where AY~60(c/o,,) and E,=(vy—Cosy)Do/e is the motional electric field in the ESW frame."," In fact, the energy gain of electron from the motional electric field $\Delta \varepsilon_m$ can be estimated as $e E_m \Delta Y$, where $\Delta Y \sim 60 (c/\omega_{pe})$ and $E_m = (v_0 - v_{\rm esw}) B_0 /c$ is the motional electric field in the ESW frame."
+" Since (he propagation speed of ESW is almost same as the speed of the reflected ions (ShimadaandIoshino2000).. we may assume (ju,μμ."," Since the propagation speed of ESW is almost same as the speed of the reflected ions \citep{Non00}, we may assume $v_{\rm esw} \sim -v_0$."
+" Then we get Xen020oufoyηem,οὗτο1.6Gm,c. and the Lorentz [actor of the accelerated electron becomes 2.6."," Then we get $\Delta \varepsilon_m \sim 120 (\omega_{ce}/\omega_{pe})(v_0/c) m_e c^2 \sim 1.6 m_e c^2$, and the Lorentz factor of the accelerated electron becomes 2.6."
+" Namely. (he normalized momentum D./12,9. which is almost consistent with ihe momentum gain in Figure 3."," Namely, the normalized momentum $P_e/P_0 \sim 9$, which is almost consistent with the momentum gain in Figure 3."
+" From the above analvsis of the simulation data. we think the so called. ""shock surfing mechanism plavs an important role on electron acceleration."," From the above analysis of the simulation data, we think the so called “shock surfing” mechanism plays an important role on electron acceleration."
+ Let us quickly review (he idea of the shock surfing (Saedeev1966:SagdeevandShapiro1973:SugiharaMizuno1979).," Let us quickly review the idea of the shock surfing \citep{Sag66,Sag73,Sug79}."
+. The shock surfing mechanism has been extensively studied for the ion acceleration 2001)..," The shock surfing mechanism has been extensively studied for the ion acceleration \citep{Kat83,Lem84,Ohs85,Zan96,Lee96,Uce01}. ."
+ Due to the inertia difference between ions aud electrons flowing into the shock. the polarization electric field normal to the shock front is formed.," Due to the inertia difference between ions and electrons flowing into the shock, the polarization electric field normal to the shock front is formed."
+ An ion having a, An ion having a
+"show that even for a point-like TNR ignition, the expansion and progress of the runaway is almost spherically symmetric for nova conditions.","show that even for a point-like TNR ignition, the expansion and progress of the runaway is almost spherically symmetric for nova conditions."
+" We note that the adopted resolution as well as the size, intensity, and location of the initial perturbation have a very limited impact on the results, principally affecting the timescale for the onset of the KH instability but not the final, mean metallicity."," We note that the adopted resolution as well as the size, intensity, and location of the initial perturbation have a very limited impact on the results, principally affecting the timescale for the onset of the KH instability but not the final, mean metallicity."
+ Details will be extensively discussed in a forthcoming publication., Details will be extensively discussed in a forthcoming publication.
+" Our results agree with earlier 2-D hydrodynamic simulations (GLT97) and solve the controversy raised by another 2-D study (KHT98) that questioned the efficiency of this mixing mechanism, and hence the corresponding strength of the runaway and its capability to power a fast nova outburst."," Our results agree with earlier 2-D hydrodynamic simulations (GLT97) and solve the controversy raised by another 2-D study (KHT98) that questioned the efficiency of this mixing mechanism, and hence the corresponding strength of the runaway and its capability to power a fast nova outburst."
+up to a normalization factor.,up to a normalization factor.
+" To specify a dynamical system consistent with this invariant distribution, we must also make a choice of one form (2."," To specify a dynamical system consistent with this invariant distribution, we must also make a choice of one form $\Omega$."
+" Consider With this choice, gives the dynamical system: Note that the equations clearly do not permit passage across y—0 or z—0."," Consider With this choice, gives the dynamical system: Note that the equations clearly do not permit passage across $y=0$ or $z=0$."
+" Thus the initial distribution, which has support for positive and negative y and z, can not be ergodic."," Thus the initial distribution, which has support for positive and negative $y$ and $z$, can not be ergodic."
+ Let us therefore consider a modified probability density up to normalization., Let us therefore consider a modified probability density up to normalization.
+" Although this modification introduces delta functions in 0,p and ὃμρ at the x=0 and y=0 boundaries respectively, the Fokker Planck equation is still satisfied because vz(a=0)0 and vy(y=0)0."," Although this modification introduces delta functions in $\partial_x\rho$ and $\partial_y \rho$ at the $x=0$ and $y=0$ boundaries respectively, the Fokker Planck equation is still satisfied because $v_x(x=0) =0$ and $v_y(y=0) =0$."
+" Figure 1 shows planar projections of points generated both by numerical simulation of this dynamical system, starting from the initial condition 0.5), and by random sampling according to the probability distribution(3."," Figure 1 shows planar projections of points generated both by numerical simulation of this dynamical system, starting from the initial condition $(x,y,z)=(0.5,0.5,0.5)$ , and by random sampling according to the probability distribution."
+"17).. At the crude level of visual inspection, these seem to agree."," At the crude level of visual inspection, these seem to agree."
+ Figure 1., Figure 1.
+" The XY and YZ projection of points generated by the dynamical system observed at fixed time intervals At=4 (on the left) starting from the initial condition (2, and by random sampling with the probability distribution (on the right)."," The XY and YZ projection of points generated by the dynamical system observed at fixed time intervals $\Delta t=4$ (on the left) starting from the initial condition $(x,y,z)=(0.5,0.5,0.5)$ and by random sampling with the probability distribution (on the right)."
+"an image separation of z1&10! aresec for τι=0.5 aud 7,=2.0.",an image separation of $\approx 4\times 10^{-4}$ arcsec for $z_\mathrm{l}=0.5$ and $z_\mathrm{s}=2.0$.
+ This is sufficient to allow VSOP-2 to resolve the image splitting. regardless of the density profile of the subhalo hosting the IAIBIL," This is sufficient to allow VSOP-2 to resolve the image splitting, regardless of the density profile of the subhalo hosting the IMBH."
+ Of course. even if the Ferrarese relatious would hold for huuious cwiut ealaxies. the do not necessarily do so for dark ones. since this depends on the formation details of IAIBTs.," Of course, even if the Ferrarese relations would hold for luminous dwarf galaxies, the do not necessarily do so for dark ones, since this depends on the formation details of IMBHs."
+ Moreover. a population of halo IXIBIsnot associated with subhalos could forma au undesired backerouud of nullilensing events that would obfuscate attempts to study subhalos through image-splitting effects.," Moreover, a population of halo IMBHs associated with subhalos could form an undesired background of millilensing events that would obfuscate attempts to study subhalos through image-splitting effects."
+ Finally. there is an imiportaut issue related to the limitation of current CDM halo simulations.," Finally, there is an important issue related to the limitation of current CDM halo simulations."
+ Cirrent N-hody simulations oulv resolve scales down to Ro-- 000. whereas for the steeper deusitv profiles ic. those producing detectable nage separations for dwartealaxy masses a non-neslieible fraction of X(«p) at the Enmstein radius comes from radii smaller hau this.," Current N-body simulations only resolve scales down to $R\sim 0.001 r_\mathrm{vir}$ , whereas for the steeper density profiles -- i.e. those producing detectable image separations for dwarf-galaxy masses – a non-negligible fraction of $\bar{\Sigma}(<r)$ at the Einstein radius comes from radii smaller than this."
+ This ialkes the exercise of deriving dark ido lensing properties from fitting functions based on current N-hody data somewhat hazardous., This makes the exercise of deriving dark halo lensing properties from fitting functions based on current N-body data somewhat hazardous.
+ Uuless heoretical arguments(c.g.theTausen&Stadel2006). can © used to determune slope of the density profile at even smaller radi robust predictions for the nuage separations produced bv low-nass halos may be very difficult to derive.," Unless theoretical arguments \citep[e.g.][]{Hansen & Stadel} can be used to determine the slope of the density profile at even smaller radii, robust predictions for the image separations produced by low-mass halos may be very difficult to derive."
+ Observational limits ou the lensing optical depth as a fiuction of A0. based on searches for subhalo image splitting. may eventually turi out to be a useful way of setting constraints ou the internal structure of dark matter halos at scales below the resolution of nuinerical simulations.," Observational limits on the lensing optical depth as a function of $\Delta\theta$, based on searches for subhalo image splitting, may eventually turn out to be a useful way of setting constraints on the internal structure of dark matter halos at scales below the resolution of numerical simulations."
+ Of course. the optical depth does not depend only on the deusitv profile of the sublialos. but also on their mass fiction aud the spatial distribution within the macrolenus.," Of course, the optical depth does not depend only on the density profile of the subhalos, but also on their mass function and the spatial distribution within the macrolens."
+ Detailed estimates of +he optical depth for stroug lensing by subhalos in tle wartealaxy mass range is outside the scope of this study. but will be covered in paper I of this series (Richi ct al.," Detailed estimates of the optical depth for strong lensing by subhalos in the dwarf-galaxy mass range is outside the scope of this study, but will be covered in paper II of this series (Riehm et al."
+ 2008. iu preparation).," 2008, in preparation)."
+within 10 AU of the star (Robitailleetal.2007).,within 10 AU of the star \citep{ro07}.
+". If, however, the inner disk is only partially depleted and dust emission remains optically thick in the infrared, the observed double intensity peak and the spectral energy distribution can be reconciled."," If, however, the inner disk is only partially depleted and dust emission remains optically thick in the infrared, the observed double intensity peak and the spectral energy distribution can be reconciled."
+ A number of physical mechanisms could reduce the dust density in the inner region of circumstellar disks., A number of physical mechanisms could reduce the dust density in the inner region of circumstellar disks.
+" For example, planets less massive than Jupiter may carve partially depleted gaps in the surface density distribution by tidal interaction with the surrounding material (Brydenetal. 1999)."," For example, planets less massive than Jupiter may carve partially depleted gaps in the surface density distribution by tidal interaction with the surrounding material \citep{br99}."
+. This possibility is discussed in more detail in Section 5.3.., This possibility is discussed in more detail in Section \ref{sec:surf}.
+" In Paper I, we also proposed that a surface density profile that gradually decreases towards the star may originate naturally from the viscous evolution of a disk if viscosity"," In Paper I, we also proposed that a surface density profile that gradually decreases towards the star may originate naturally from the viscous evolution of a disk if viscosity"
+The spread of metallicity in the thin disk at the solar radius is mostly due to non-evolutionarv effects: the increase in metallicity due to chemical evolution eveling over the last 8 Gvr (the age of the thin disk according to solar neighborhood data. see e.g Lin Chabover 2000) is limited to about [Fe/II]--0.4 (from ~ -0.2 to 70.2 dex). while the range of measured metallicity is of the order of 1.2 dex (from ~-0.6 to ~0.6 dex).,"The spread of metallicity in the thin disk at the solar radius is mostly due to non-evolutionary effects: the increase in metallicity due to chemical evolution cycling over the last 8 Gyr (the age of the thin disk according to solar neighborhood data, see e.g Liu Chaboyer 2000) is limited to about $\sim$ 0.4 (from $\sim$ -0.2 to $\sim$ 0.2 dex), while the range of measured metallicity is of the order of 1.2 dex (from $\sim$ -0.6 to $\sim$ 0.6 dex)."
+ How these eXtreme metallicities are compatible with the gradients measured in the previous section?, How these extreme metallicities are compatible with the gradients measured in the previous section?
+ The reason for this spread is to be found in stars that migrate across the disk (Haywood 2003)., The reason for this spread is to be found in stars that migrate across the disk (Haywood 2008).
+ Suggestions and results by Grenon (1972). Sellwood Binney (2002). and. Haywood (2008) conclusively point to the fact that local metal-rich (|Fe/11]2-0.2 dex) and metal-poor «-0.2 dex) stars of the thin disk have been brought to the solar radius by radial migration.," Suggestions and results by Grenon (1972), Sellwood Binney (2002), and Haywood (2008) conclusively point to the fact that local metal-rich $>$ 0.2 dex) and metal-poor $<$ -0.2 dex) stars of the thin disk have been brought to the solar radius by radial migration."
+ Sellwood Binnev (2002) proposed a mechanism (hat allows stars to individually exchange (heir angular momentum. allowing them (ο swap from one nearly circular orbit to another of different mean orbital radius.," Sellwood Binney (2002) proposed a mechanism that allows stars to individually exchange their angular momentum, allowing them to swap from one nearly circular orbit to another of different mean orbital radius."
+ At the solar radius. (his mechanism involves a moderate fraction of stars. but itis important since il eives a glimpse ol the inner and outer Galactic disk chemistry.," At the solar radius, this mechanism involves a moderate fraction of stars, but it is important since it gives a glimpse of the inner and outer Galactic disk chemistry."
+ The highest and lowest thin disk iron abundances measured in the solar neighborhood are respectively 0.6 and ~-0.6 dex. corresponding to [O/11]9-0.3 and -0.3 dex (Ramurrez οἱ al.," The highest and lowest thin disk iron abundances measured in the solar neighborhood are respectively $\sim$ 0.6 and $\sim$ -0.6 dex, corresponding to $\sim$ 0.3 and -0.3 dex (Ramírrez et al."
+ 2007. Fig.," 2007, Fig."
+ 8). which roughly corresponds to the factor of 1/2 that is often used when scaling iron to oxvgen abundances (e.g Martinelli Matteucci 2000. or more recently. Caimnmi 2008).," 8), which roughly corresponds to the factor of 1/2 that is often used when scaling iron to oxygen abundances (e.g Martinelli Matteucci 2000, or more recently, Caimmi 2008)."
+ The radial oxvgen gradient for intermediate age PNe (Type 1). ~-0.023 dex |. would correspond to [Fe/II]--0.05Hr dex ! if using the conversion factor illustrated above.," The radial oxygen gradient for intermediate age PNe (Type II), $\sim$ -0.023 dex $^{-1}$, would correspond to $\sim$ -0.05 dex $^{-1}$ if using the conversion factor illustrated above."
+ Given that the solar neighborhood mean metallicity is 20.0 dex. such radial gradient," Given that the solar neighborhood mean metallicity is $\sim$ 0.0 dex, such radial gradient"
+"a similar procedure. bul using spectra of emission-line galaxies rather (han models. ? uses five photometric bands (D. Re. V. / and 2"").","a similar procedure, but using spectra of emission-line galaxies rather than models, \citet{2006astro.ph.10846L} uses five photometric bands $B$ , $R_C$ , $V$, $i'$ and $z'$ )."
+ survevs use dropout technique in fillers bluer than 7 in order to assure that the Lyman break (A = 912A)) lies bluewards the narrow-band filter., surveys use dropout technique in filters bluer than $i$ in order to assure that the Lyman break $\lambda$ = ) lies bluewards the narrow-band filter.
+ For example. ? request that the eandidates are not detected in B. V. and HR. in addition to other criteria color criteria.," For example, \citet{2005PASJ...57..165T}
+ request that the candidates are not detected in $B$, $V$ and $R$, in addition to other criteria color criteria."
+ Finally. we mention a more complex classification scheme.," Finally, we mention a more complex classification scheme."
+ It is presented in ? and il is used in the CADIS survey (?).., It is presented in \citet{2001A&A...365..681W} and it is used in the CADIS survey \citep{2003A&A...402...65H}.
+ Twelve medium pass band fillers are used. in addition to broad-band D. H. £. J ancl WN. to classify the objects in three broad. categories: stars. quasars (including AGN) and galaxies.," Twelve medium pass band filters are used, in addition to broad-band $B$, $R$, $I$, $J$ and $K'$, to classify the objects in three broad categories: stars, quasars (including AGN) and galaxies."
+ Fluxes of the objects through the filters are obtained integrating the spectrum of the object through the filter profile., Fluxes of the objects through the filters are obtained integrating the spectrum of the object through the filter profile.
+ The narrow-bancl filter technique allows. not only. (o select 1e objects wilh a possible emission line. but also to compute the value of the line flux aud je equivalent width of the line.," The narrow-band filter technique allows, not only to select the objects with a possible emission line, but also to compute the value of the line flux and the equivalent width of the line."
+ In order to obtain the line and continuum fIux for (he selected objects. several simplifications. iab are studied in this section. need to be made.," In order to obtain the line and continuum flux for the selected objects, several simplifications, that are studied in this section, need to be made."
+ First. the line profile is assumed infinitely in when compared with the width of the πατον)Ρα. filler.," First, the line profile is assumed infinitely thin when compared with the width of the narrow-band filter."
+ Secoucl. (he positioning of je emission within the transmittance profile of the used fillers is very close to the center of je narrow-band filler.," Second, the positioning of the emission within the transmittance profile of the used filters is very close to the center of the narrow-band filter."
+ With simple assumption we can obtain a mean wavelength lor the enission line., With simple assumption we can obtain a mean wavelength for the emission line.
+ Finally. the continuum flux can be modeled by an analvtic function. either a power-law or a polynomial.," Finally, the continuum flux can be modeled by an analytic function, either a power-law or a polynomial."
+ It is critical to have a good estimate of the continuum flux., It is critical to have a good estimate of the continuum flux.
+ The FWIIM of the lines of star-Forming galaxies is related with the mass of the object and typically is less than ~ 10A., The FWHM of the lines of star-forming galaxies is related with the mass of the object and typically is less than $\sim$ 10.
+. To estimate the inlluence of the finite width of the emission line we have integrated Gaussian profiles through. Gaussian and rectangular narrow-band filters ancl computed the recovered line flux., To estimate the influence of the finite width of the emission line we have integrated Gaussian profiles through Gaussian and rectangular narrow-band filters and computed the recovered line flux.
+ The EWILIM increases with redshift a factor 1Ἔαν πο the infinitely thin line approximation can be valid. recovering more than of the line flix. up to ze4 lor very narrow filters (A~ 50:1).," The FWHM increases with redshift a factor $1+z$, so the infinitely thin line approximation can be valid, recovering more than of the line flux, up to $z\sim4$ for very narrow filters $\Delta\sim50\AA$ )."
+ For wider filters. (A~ 150-4). we recover more than of the flux up to ze10.," For wider filters, $\Delta\sim150\AA$ ), we recover more than of the flux up to $z\sim10$."
+ These limit redshifts awe lower [or sources with broader lines. such as AGN.," These limit redshifts are lower for sources with broader lines, such as AGN."
+ The ELGscan stillbe selected using the methods described in Sect., The ELGscan stillbe selected using the methods described in Sect.
+5- but the fluxes of the lines computed with this section's equations will,\ref{sec:dispq} but the fluxes of the lines computed with this section's equations will
+ie. R=τα:+y). where 2 is the forbidden line and «+y represents the sum of the infercombination emission.,"i.e. $R = z/(x+y)$, where $z$ is the forbidden line and $x+y$ represents the sum of the intercombination emission."
+ In (he 4445 spectrum (he best constraints on J? arises from ihe triplet., In the 445 spectrum the best constraints on $R$ arises from the triplet.
+" Allowing the ratio of the forbidden and intercombination lines (o vary results in a value of 2=0.9m (al confidence). which implies an electron densitv in the region of n,=LOM—10! cem.? (e.g. Figure 8. Porquet&Dubau2000))."," Allowing the ratio of the forbidden and intercombination lines to vary results in a value of $R=0.9^{+1.1}_{-0.3}$ (at confidence), which implies an electron density in the region of $n_{\rm e}=10^{10}-10^{11}$ $^{-3}$ (e.g. Figure 8, \citealt{porquet00}) )."
+ This suggests (he emission originates [from matter closer in (han the Narrow Line Region (NLR). a point we discuss further in Section 5.," This suggests the emission originates from matter closer in than the Narrow Line Region (NLR), a point we discuss further in Section 5."
+ The velocity widths of the ancl emission lines were also determined., The velocity widths of the and emission lines were also determined.
+ The Lyviman-a line is resolved with a width of o=2.0H eeV (see Table 1). corresponding to a FWHM width of 2100.12km ss.|: the fit statistic worsens by AC’=12 if the line width is set to zero.," The $\alpha$ line is resolved with a width of $\sigma=2.0^{+1.1}_{-0.9}$ eV (see Table 1), corresponding to a FWHM width of $2100^{+1100}_{-950}$ $^{-1}$; the fit statistic worsens by $\Delta C=12$ if the line width is set to zero."
+ Note that the separation of the Lyman-a doublet has a negligible effect on the line width., Note that the separation of the $\alpha$ doublet has a negligible effect on the line width.
+ Upon modeling the ]le-like triplet emission using (wo separate Gaussian lines (o represent the forbidden ancl intercombinalion enission. the He-o line widthwas also determined. giving 6=2.4m eeV (see Table 1). which corresponds to a FWIIM. width of 300024 |.," Upon modeling the He-like triplet emission using two separate Gaussian lines to represent the forbidden and intercombination emission, the $\alpha$ line widthwas also determined, giving $\sigma=2.4^{+2.8}_{-0.8}$ eV (see Table 1), which corresponds to a FWHM width of $3000^{+3400}_{-980}$ $^{-1}$."
+ Note that the line widths of the forbidden ancl intercombination lines were assumed to be equal to each other in the model., Note that the line widths of the forbidden and intercombination lines were assumed to be equal to each other in the model.
+ The fit statistic also worsened significantly (AC= +31) when the velocity width of the lines were fixed to zero., The fit statistic also worsened significantly $\Delta C=+31$ ) when the velocity width of the lines were fixed to zero.
+ ILowever the lines from higher Z ions were not resolved. as the LETC spectral resolution worsens with increasing energv. however (he upper-limits to their widths are consistent with the values from the Oxygen lines (see Table 1 for the II-Iike ions and Table 2 for the He-like ions).," However the lines from higher Z ions were not resolved, as the LETG spectral resolution worsens with increasing energy, however the upper-limits to their widths are consistent with the values from the Oxygen lines (see Table 1 for the H-like ions and Table 2 for the He-like ions)."
+ In order to derive the most accurate determination of the line velocity width. we assumed that the widths of the three strongest lines [rom lexiscviitt£). Lexisevii((i) and Lvman-a were identical and tied (hese values in the resulting model.," In order to derive the most accurate determination of the line velocity width, we assumed that the widths of the three strongest lines from (f), (i) and $\alpha$ were identical and tied these values in the resulting model."
+ The resolution ol the LETG is also at its highest in the O band., The resolution of the LETG is also at its highest in the O band.
+" This vielded a best-fit velocity width of c=1120Zo ! (or e=2.1PS eeV al 56leeV). corresponding to a FWIIM width of Myyuy,=2600.DP |..."," This yielded a best-fit velocity width of $\sigma=1120^{+430}_{-270}$ $^{-1}$ (or $\sigma=2.1^{+0.8}_{-0.5}$ eV at eV), corresponding to a FWHM width of $v_{\rm FWHM}=2600^{+1000}_{-600}$ $^{-1}$ ."
+ A contour plot showing the measurement of the FWIIM, A contour plot showing the measurement of the FWHM
+Since the mass of the black hole in nearby galaxies appears to be proportional to the spheroid mass. the mass function of black holes must be similar in shape to the spheroid mass fiction.,"Since the mass of the black hole in nearby galaxies appears to be proportional to the spheroid mass, the mass function of black holes must be similar in shape to the spheroid mass function."
+ The mean black hole mass is therefore that appropriate to au LS galaxy. or about 3«105 AL...," The mean black hole mass is therefore that appropriate to an $L^*$ galaxy, or about $3\times 10^8$ $_{\odot}$."
+ The Eddingtou limit of such a black hole is about 3«10! eye 1 and its mass doubling (Salpeter) time is about ὃνLOT vy., The Eddington limit of such a black hole is about $3\times 10^{46}$ erg $^{-1}$ and its mass doubling (Salpeter) time is about $3\times 10^7$ yr.
+ If he typical mass black hole has therefore grown from sav a 1aillion solar mass one n3«10 (ie by :— 2). then sve probably need L>0.405Zgj;19 ere s," If the typical mass black hole has therefore grown from say a million solar mass one in $3\times
+10^9$ yr (i.e. by $z\sim 2$ ), then we probably need $L>0.05
+L_{Edd}\sim 10^{45}$ erg $^{-1}$."
+ This means that the typical erowing black hole was powerful aud of quasar-like iuninositv (indeed housing a quasar at the centre)., This means that the typical growing black hole was powerful and of quasar-like luminosity (indeed housing a quasar at the centre).
+ Such an obscured powerful object would locally be classified as à ULIRG (see Sanders Mirabol 1996). although the distaut ones need not be the same as the ocal ones. which are perhaps mainly fuclled by mergers.," Such an obscured powerful object would locally be classified as a ULIRG (see Sanders Mirabel 1996), although the distant ones need not be the same as the local ones, which are perhaps mainly fuelled by mergers."
+ Of course it is possible that massive black holes erew inside galaxies which hemsclves were iiereiug back at 22., Of course it is possible that massive black holes grew inside galaxies which themselves were merging back at $z\sim 2$.
+ Nevertheless. unless they were all assembled roni sinaller holes just before accretion switched off. it is probable that they cuit or a reasonable fraction of the last doubling time as a single object.," Nevertheless, unless they were all assembled from smaller holes just before accretion switched off, it is probable that they emit for a reasonable fraction of the last doubling time as a single object."
+ Consider an isothermal galaxy im which a significant fraction f. of cooled gas reais as cold dusty clouds instead of rapidly forming stars., Consider an isothermal galaxy in which a significant fraction $f_c$ of cooled gas remains as cold dusty clouds instead of rapidly forming stars.
+ At the centre a black hole erows by acerction from the surrounding cold (and hot) eas., At the centre a black hole grows by accretion from the surrounding cold (and hot) gas.
+" Asstmue that the nucleus also blows a wind of velocity ο which has a power L,=aLeg.", Assume that the nucleus also blows a wind of velocity $v_w$ which has a power $L_w=\alpha L_{Edd}$.
+ Eveutually the wiud becomes powerful chough to blow away the surrounding gas aud so shut off the accretion and further erowth to the black hole aud spheroid., Eventually the wind becomes powerful enough to blow away the surrounding gas and so shut off the accretion and further growth to the black hole and spheroid.
+ The Magorrian ct al (1998) black-hole spheroid mass relation can then be obtained (Silk Ress 1998: Fabian 1999: Dlandford 1999)., The Magorrian et al (1998) black-hole – spheroid mass relation can then be obtained (Silk Ress 1998; Fabian 1999; Blandford 1999).
+" The kinetic power of a wind at which it ejects cold gas of column density Vy, from a spheroid is given by Or where 7 is the velocity dispersion within the spheroid. (", The kinetic power of a wind at which it ejects cold gas of column density $N_H$ from a spheroid is given by or where $\sigma$ is the velocity dispersion within the spheroid. (
+"I have used a force argument here. see Fabian 1999: Silk Rees 1998 use au cherey aremuecnt to obtain a Imt of o/G. which is a factor 0/0, smaller than the above £,,..)","I have used a force argument here, see Fabian 1999; Silk Rees 1998 use an energy argument to obtain a limit of $\sigma^5/G$, which is a factor $\sigma/v_w$ smaller than the above $L_w$ .)"
+ Ejection occurs when Using the Faber-Jacksou relation for sphoroids (Mi;x7 1) thon vields. if »L—] ," Ejection occurs when Using the Faber-Jackson relation for spheroids $M_{sph}\propto \sigma^4$ ) then yields, if ${v_w\over c}{f_c\over \alpha}\sim 1$ "
+"satellite galaxies aside, the studies ofMarin-Franchetal.(2009),, Dotteretal.(2010),, and Dotteretal.(2011) show that the age-metallicity relation of the GGCs possesses two branches — one with near constant old age of ~13 Gyrs and another branch to significantly younger ages with increased [Fe/H].","satellite galaxies aside, the studies of\citet{Marin-Franch09}, , \citet{Dotter10}, , and \citet{Dotter11} show that the age-metallicity relation of the GGCs possesses two branches – one with near constant old age of $\sim 13$ Gyrs and another branch to significantly younger ages with increased [Fe/H]."
+ The younger branch of GGCs is found to be dominated by GCs associated with the Sagittarius and Canis Major dwarf galaxies., The younger branch of GGCs is found to be dominated by GCs associated with the Sagittarius and Canis Major dwarf galaxies.
+ This leads to conclude that approximately one quarter of the GGC system has resulted from accretion., This leads \citet{Forbes10} to conclude that approximately one quarter of the GGC system has resulted from accretion.
+" In this study, we show that the YH GCs of the MW (as defined by a combination of age estimators) are confined to a plane that is indistinguishable from the MW PoS. The YH GCs act as tracers of the dwarf galaxies that have been disrupted in the formation of the outer halo."," In this study, we show that the YH GCs of the MW (as defined by a combination of age estimators) are confined to a plane that is indistinguishable from the MW PoS. The YH GCs act as tracers of the dwarf galaxies that have been disrupted in the formation of the outer halo."
+ The consistency in spatial alignment of early accretion events to the alignment of current satellites suggests a preferred orientation for accretion from the local large scale structure that has remained consistent for a large fraction of a Hubble time., The consistency in spatial alignment of early accretion events to the alignment of current satellites suggests a preferred orientation for accretion from the local large scale structure that has remained consistent for a large fraction of a Hubble time.
+" When excellent self-consistent photometry exists, such as in the study of (2009),, it is possible to derive accurate relative ages even amongst clusters of advanced age."," When excellent self-consistent photometry exists, such as in the study of \citet{Marin-Franch09}, it is possible to derive accurate relative ages even amongst clusters of advanced age."
+ These authors compare the luminosity of the main-sequence turn-off to that of the lower main-sequence and hence express the relative age for a sample of 64 nearby GCs., These authors compare the luminosity of the main-sequence turn-off to that of the lower main-sequence and hence express the relative age for a sample of 64 nearby GCs.
+" finds evidence for two groups of GCs: an old GC population with an age dispersion of ~5% and no age-metallicity relation, and a young GC population with an age-metallicity relation similar to that defined by the GCs associated with the Sagittarius dwarf galaxy."," \citeauthor{Marin-Franch09} finds evidence for two groups of GCs: an old GC population with an age dispersion of $\sim5$ and no age-metallicity relation, and a young GC population with an age-metallicity relation similar to that defined by the GCs associated with the Sagittarius dwarf galaxy."
+" The age determination method of Zinn(1993) relies on a relation between age, metallicity, and HB morphology to resolve relative cluster age differences."," The age determination method of \citet{Zinn93} relies on a relation between age, metallicity, and HB morphology to resolve relative cluster age differences."
+ Zinn defined a fiducial sequence for inner halo clusters (Hg.<6 kpc) in a against [Fe/H] diagram and then measured the difference between the HB-type of a given cluster and that of the fiducial sequence., \citeauthor{Zinn93} defined a fiducial sequence for inner halo clusters $R_{gc} < 6$ kpc) in a against [Fe/H] diagram and then measured the difference between the HB-type of a given cluster and that of the fiducial sequence.
+" Those with a difference in HB-type > —0.4 were classified by Zinn(1993) as young clusters, the remainder old clusters."," Those with a difference in HB-type $>$ $-0.4$ were classified by \citet{Zinn93} as young clusters, the remainder old clusters."
+ Mackey&Gilmore(2004) has revisited the issue of relative age partition., \citet{Mackey04} has revisited the issue of relative age partition.
+ Rather than using spatial information in the derivation of their fiducial these authors utilise the model evolutionary isochrones of Reyetal. (2001)., Rather than using spatial information in the derivation of their fiducial these authors utilise the model evolutionary isochrones of \citet{Rey01}. .
+ Theisochronecorresponding to the canonical ancient GC population traces the those clusters, Theisochronecorresponding to the canonical ancient GC population traces the those clusters
+between the spectroscopically inferred T.spectro and the true Teg amounts to up to5%.,between the spectroscopically inferred $\teffs$ and the true $\teff$ amounts to up to.
+". To see how the results may be affected by stellar activity, we first compare standard evolution models to evolution models calculated with CESAM and a modified atmospheric boundary condition Figure 1 shows the result for a given luminosity and age, a star with spots simply has a higher ""spectroscopic"" effective temperature Teg,spectro aS inferred from spectroscopic measurements than a star without spots."," To see how the results may be affected by stellar activity, we first compare standard evolution models to evolution models calculated with CESAM and a modified atmospheric boundary condition Figure \ref{fig:hr_comp} shows the result for a given luminosity and age, a star with spots simply has a higher “spectroscopic” effective temperature $\teffs$ as inferred from spectroscopic measurements than a star without spots."
+ The evolution is simply displaced to the left of the HR diagram by analmost constant ratio (1—y;)!/4 in Teg., The evolution is simply displaced to the left of the HR diagram by analmost constant ratio $(1-\chi_{\rm s})^{1/4}$ in $\teff$.
+" Quantitatively, the mean deviation amounts to 5x10? on Tey, with a standard deviation o=4.6x107T,q and a maximum deviation 6.7x10:2Τ./."," Quantitatively, the mean deviation amounts to $5\times 10^{-5}$ on $\modif \teff$, with a standard deviation $\sigma=4.6\times 10^{-4}\modif\teff$ and a maximum deviation $6.7\times 10^{-3}\modif\teff$."
+" The small departures from this constant displacement are due to slight modifications of the atmospheric properties (opacities) with temperature, but they can safely be neglected in this study."," The small departures from this constant displacement are due to slight modifications of the atmospheric properties (opacities) with temperature, but they can safely be neglected in this study."
+" When considering a star's evolution, that a star has spots is therefore equivalent to an added uncertainty in the measurement of its effective temperature."," When considering a star's evolution, that a star has spots is therefore equivalent to an added uncertainty in the measurement of its effective temperature."
+ A larger error bar in Tegused as a proxyfor the activity and starspots., A larger error bar in $\teff$ as a proxyfor the activity and starspots.
+" The star's physical parameters (M,, R,, age) are obtained by matching the constraints from Table | to a grid of evolution models, as depicted in Fig. 2.."," The star's physical parameters $M_\star$, $R_\star$, age) are obtained by matching the constraints from Table \ref{tab:obs} to a grid of evolution models, as depicted in Fig. \ref{fig:evolution}."
+ The two most important constraints of the problem are the star’s effective temperature and density., The two most important constraints of the problem are the star's effective temperature and density.
+" A third constraint (not shown on the plot) is the star’spresent metallicity, which should be compared to the one obtained from the evolution models that include diffusion."," A third constraint (not shown on the plot) is the star's metallicity, which should be compared to the one obtained from the evolution models that include diffusion."
+" One could include other constraints (such as that on logg), but in the present case, they are too weak to be useful."," One could include other constraints (such as that on $\log g$ ), but in the present case, they are too weak to be useful."
+" The quality of the fit of any given model is given by its distance n,, to the ellipsoid of constraints, measured in units of the standard error in the constraints given by where X; arethe constraints and o;; their standard deviation (assumed Gaussian)."," The quality of the fit of any given model is given by its distance $n_{\sigma_\star}$ to the ellipsoid of constraints, measured in units of the standard error in the constraints given by where $X_i$ arethe constraints and $\sigma_i$ their standard deviation (assumed Gaussian)."
+" The ellipsoid of constraints (of dimension 2) is centered on (T.g,p,) and has semi-minor and semi- axes Kop,Πσίστιι,05,), respectively, where Kop, is the quantile of a 2D gaussian law at the equivalent level of confidence n,."," The ellipsoid of constraints (of dimension 2) is centered on $\teff,\rho_\star$ ) and has semi-minor and semi-major axes $k_{\rm 2D,n_{\sigma}} n_{\sigma}
+(\sigma_{\teff},\sigma_{\rho_\star})$, respectively, where $k_{\rm
+ 2D,n_{\sigma}}$ is the quantile of a 2D gaussian law at the equivalent level of confidence $n_{\sigma}$."
+" In addition, kopa,~1.52,2.49,3.44 for no,=1,2,3 respectively."," In addition, $k_{\rm 2D, n_{\sigma}}
+\sim 1.52,\, 2.49,\, 3.44$ for $n_{\sigma} = 1,2,3$ respectively."
+" This normalization ensures that our solutions at 1,2,3 σ have the correct probability of occurrence."," This normalization ensures that our solutions at 1,2,3 $\sigma$ have the correct probability of occurrence."
+" However, for the metallicity we adopt a relatively crude simplification: we consider as valid only solutions for which the [Fe/H] value is within Ίσῃεε/ηι of the measured one."," However, for the metallicity we adopt a relatively crude simplification: we consider as valid only solutions for which the $\rm[Fe/H]$ value is within $\sigma_{\rm[Fe/H]}$ of the measured one."
+" We tested that in the particular case of CoRoT-2, considering yet larger errors in [Fe/H] has a negligible effect."," We tested that in the particular case of CoRoT-2, considering yet larger errors in $\rm[Fe/H]$ has a negligible effect."
+" In the remainder of the paper, we present models for which ng,<1,2,3, corresponding to confidence levels of68.396,95.4%,, and 99.7%,, respectively."," In the remainder of the paper, we present models for which $n_{\sigma_\star}\le 1,2,3$, corresponding to confidence levels of, and , respectively."
+" Two possible values of the effective temperature are used: (i) in the no-spot case, we assume that obtained by ? but with a slightly larger error to account for possible systematic errors (Tem,spectro=5608+80 KK); (ii) when including the effect of spots, we define a new temperature and its associated error Ar, to take into account the presence of spots (from to of the areaof the star) We thus derive Teg=5224—5688 KK at lo, Teg KK at 2c, and Teg=5064—5848 KK at 3c."," Two possible values of the effective temperature are used: (i) in the no-spot case, we assume that obtained by \citet{AmmlervonEiff+09} but with a slightly larger error to account for possible systematic errors $\teffs=5608\pm 80$ K); (ii) when including the effect of spots, we define a new temperature and its associated error $\Delta_{\teff}$ to take into account the presence of spots (from to of the areaof the star) We thus derive $\teff = 5224-5688$ K at $\sigma$ , $\teff=5144-5768$ K at $\sigma$ , and $\teff=5064-5848$ K at $\sigma$ ."
+ The errors are thus intrinsically non-Gaussian., The errors are thus intrinsically non-Gaussian.
+" The constraints used for the stellar density and metallicity are those derivedby ? and ?,, respectively (see Table 1))."," The constraints used for the stellar density and metallicity are those derivedby \citet{Gillon+10}
+ and \citet{Alonso+08}, , respectively (see Table \ref{tab:obs}) )."
+Present ideas on the origin of SO galaxies are mainly based on the early investigations bv Spitzer Daade (1951). Gunn Gott (1972) ancl Moore et al. (,"Present ideas on the origin of S0 galaxies are mainly based on the early investigations by Spitzer Baade (1951), Gunn Gott (1972) and Moore et al. ("
+1996).,1996).
+ Furthermore (Dekki et al., Furthermore (Bekki et al.
+ 2002) hydrodynamical interactions between spiral galaxies ancl hot intra-clIuster eas might result in gas starvation for cluster spiral galaxies., 2002) hydrodynamical interactions between spiral galaxies and hot intra-cluster gas might result in gas starvation for cluster spiral galaxies.
+ However. these papers do not explain (wo important observations: (1) Significant numbers of SO galaxies exist in the field. where ranm-pressure stripping. contact with hot intra-cluster gas and harassment should be unmmportant. and (2) The Iuminosity. distribution of SO galaxies (see Table 1) shows no significant differences between the cluster. 2roup and field environments.," However, these papers do not explain two important observations: (1) Significant numbers of S0 galaxies exist in the field, where ram-pressure stripping, contact with hot intra-cluster gas and harassment should be unimportant, and (2) The luminosity distribution of S0 galaxies (see Table 1) shows no significant differences between the cluster, group and field environments."
+ Furthermore.(3) ranrpressure stripping should be less effective in massive luminous S0 galaxies (han in [unl less massive ones.," Furthermore,(3) ram-pressure stripping should be less effective in massive luminous S0 galaxies than in faint less massive ones."
+ ILowever. the data in Table 1 appear to show no statistically significant dependence of the luminosity distribution of lenticular galaxies on environment.," However, the data in Table 1 appear to show no statistically significant dependence of the luminosity distribution of lenticular galaxies on environment."
+ Finally (4). since disk dominated Iate-tvpe galaxies predominate in the field. whereas bulge-dominate earlv-0tvpe galaxies are most common in clusters. one would have expected SO galaxies in the field (ο. on average. be flatter than those in clusters.," Finally (4), since disk dominated late-type galaxies predominate in the field, whereas bulge-dominate early-type galaxies are most common in clusters, one would have expected S0 galaxies in the field to, on average, be flatter than those in clusters."
+ However. contrary {ο expectations. the data in Table 2 shows that this is not observed to be the case.," However, contrary to expectations, the data in Table 2 shows that this is not observed to be the case."
+ This suggests that internal factors may. have been important in the transformation of some spiral galaxies into lenticulars., This suggests that internal factors may have been important in the transformation of some spiral galaxies into lenticulars.
+ The distribution of flattening values of S0 galaxies is found to be independent οἱ environment., The distribution of flattening values of S0 galaxies is found to be independent of environment.
+ This result is surprising because one might have imagined (he progenitors of SO galaxies in clusters to have mainly been early-twpe (Sa) spirals. whereas it would have been expected (hat the majority of the ancestors of SO galaxies in (he field were late-tvpe (Se) spirals.," This result is surprising because one might have imagined the progenitors of S0 galaxies in clusters to have mainly been early-type (Sa) spirals, whereas it would have been expected that the majority of the ancestors of S0 galaxies in the field were late-type (Sc) spirals."
+ Secondly. since massive luminous galaxies have deeper potential wells (han," Secondly, since massive luminous galaxies have deeper potential wells than"
+"For about 20 vears the problem of properties of short-term chauges of solar activity has been considered extensively,",For about 20 years the problem of properties of short-term changes of solar activity has been considered extensively.
+ Many investigators studied the short-term. periodicities of the various iudices of solar activity., Many investigators studied the short-term periodicities of the various indices of solar activity.
+ Several periodicities were detected. but the periodicitics about 155 days aud from the interval of |170.620] days ([1.3.1.7]. vears) are mentioned most often.," Several periodicities were detected, but the periodicities about 155 days and from the interval of $[470, 620]$ days $[1.3, 1.7]$ years) are mentioned most often."
+ Fist of them was discovered by Biegeretal.(1981) in the occurence rate of ezuunuverayv flares detected by the σααν spectrometer aboard the(SALAL)., First of them was discovered by \citet{rieg} in the occurence rate of gamma-ray flares detected by the gamma-ray spectrometer aboard the.
+. This periodicity was coufirmed for other solar fares data and for the same time period (Bogart&Bai1985:Sturrock1987:IileCliver 1991).," This periodicity was confirmed for other solar flares data and for the same time period \citep{bog, bai7, kile}."
+. It was also found iu proton flares durius solar cveles 19 and 20 (Bai&Cliver1990).. but it was uot found iu the solar flares data during solar evcles 22 (vile&Cliver1991:Bai1992:ÓzgücAtac 1989).," It was also found in proton flares during solar cycles 19 and 20 \citep{baic}, but it was not found in the solar flares data during solar cycles 22 \citep{kile, bai2, ozg}."
+. Several autors coufirmed above results for the daily sunspot area data., Several autors confirmed above results for the daily sunspot area data.
+ Lean(1990) studied the suuspot data from 18711981., \citet{lea} studied the sunspot data from 1874–1984.
+ She found the 155-day periodicity in data records from 31 vears., She found the 155-day periodicity in data records from 31 years.
+ This periodicitv is always characteristic for one of the solar hemispheres (the southern hemisphere for eveles 1215 aud the northern henüsphere for cycles 1621)., This periodicity is always characteristic for one of the solar hemispheres (the southern hemisphere for cycles 12–15 and the northern hemisphere for cycles 16–21).
+ Moreover. it is ouly present durus epochs of maxiuun activity (in episodes of 13 vears).," Moreover, it is only present during epochs of maximum activity (in episodes of 1–3 years)."
+ by Carbonell&Ballester(1992)., by \citet{car2}.
+.. They applied the sale power spectrin method as Lean. but the daily sunspot area data (eveles 1221) were divided iuto 10 shorter time series.," They applied the same power spectrum method as Lean, but the daily sunspot area data (cycles 12–21) were divided into 10 shorter time series."
+ The periocicities were searched for he frequency interval 57115 ulIz (100.200 days) aud or each of LO time senes., The periodicities were searched for the frequency interval 57–115 nHz (100–200 days) and for each of 10 time series.
+ The authors showed that he periodicity between 150160 davs is statistically siguificaut during all eveles from 16 to 21., The authors showed that the periodicity between 150–160 days is statistically significant during all cycles from 16 to 21.
+ The considered )eaks were remained unaltered after removiug the 11-vear cycle and applying the power spectrum analysis., The considered peaks were remained unaltered after removing the 11-year cycle and applying the power spectrum analysis.
+ Oliver.Ballester&Baucin(1998) used the wavelet echuique for the daily suuspot areas between 187 Laud 1993., \citet{ol8} used the wavelet technique for the daily sunspot areas between 1874 and 1993.
+ They determined the epochs of appearance of this, They determined the epochs of appearance of this
+where £j;=e);1/38;Vu and e;; is the rate of strain tensor which is given by From the dissipation we can find the torque on the disc due to the tidal perturbation.,where $E_{ij}=e_{ij}-1/3\delta_{ij} \bm{\nabla.u}$ and $e_{ij}$ is the rate of strain tensor which is given by From the dissipation we can find the torque on the disc due to the tidal perturbation.
+ The rate of working per unit surface area of the disc by a torque. 7. is where Of=dOfdr.," The rate of working per unit surface area of the disc by a torque, $T$ , is where $\Omega'=d\Omega/dr$."
+ Ehe first term is the rate of convection of energy over the whole disc., The first term is the rate of convection of energy over the whole disc.
+ Hs value depends only on the boundary conditions., Its value depends only on the boundary conditions.
+ We find the torque on the disc to be where ;/= 0. 1. 2 (Frank.Wing&Raine2002).," We find the torque on the disc to be where $i=0$ , $1$, $2$ \citep*{frank02}."
+. Phe torque Zi is defined as the torque on an accretion disc without a companion., The torque $T_0$ is defined as the torque on an accretion disc without a companion.
+ This has à viscous torque of ancl so the dissipation is In the left panel of Fig., This has a viscous torque of and so the dissipation is In the left panel of Fig.
+" 2. we plot the three scaled. dissipations. Doflv). Dj,/(CX) and Da(9X) as functions of the racius in the disc."," \ref{diss} we plot the three scaled dissipations, $D_0/(\nu\Sigma)$, $D_1/(\zeta\Sigma)$ and $D_2/(\nu\Sigma)$ as functions of the radius in the disc."
+ This is similar to Fig., This is similar to Fig.
+ 1 in Papaloizou&Pringle(1977).. but we use the Lill approximation.," 1 in \cite{papaloizou77}, but we use the Hill approximation."
+ We see that the magnitude of the dissipation from the internal viscous torques. {ο(νο). and the tidal dissipation. D»(λος are equal very close to. but just inside. the radius in the disc where the particle orbits begin to cross.," We see that the magnitude of the dissipation from the internal viscous torques, $D_0/(\nu\Sigma)$, and the tidal dissipation, $D_2/(\nu\Sigma)$, are equal very close to, but just inside, the radius in the disc where the particle orbits begin to cross."
+ In the right panel of Fig., In the right panel of Fig.
+ 2. we plot the scaled torques on the disc., \ref{diss} we plot the scaled torques on the disc.
+ Phe tidal torque in the dise starts to dominate the viscous torque just inside of the radius where the particle orbits cross., The tidal torque in the disc starts to dominate the viscous torque just inside of the radius where the particle orbits cross.
+ Llowever. outside of the radius where the particle orbits cross. pressure and nonlinear cllects cannot. be neglected.," However, outside of the radius where the particle orbits cross, pressure and nonlinear effects cannot be neglected."
+ Therefore. this linear solution is not valid bevond the orbit crossing radius.," Therefore, this linear solution is not valid beyond the orbit crossing radius."
+ We use a one-dimensional (radius only) model of an aceretion disce subject to a tidal torque in order to determine the surface density evolution., We use a one-dimensional (radius only) model of an accretion disc subject to a tidal torque in order to determine the surface density evolution.
+The governing equation for the surface density of a flat aceretion disc centered on the planet is (Pringle1981).. where --d/dr and the angular velocity is given in equation (24)).,"The governing equation for the surface density of a flat accretion disc centered on the planet is \citep{pringle81}, where $'=d/d r$ and the angular velocity is given in equation \ref{angvel}) )."
+ In the previous section. we found the tidal torque on the disc.," In the previous section, we found the tidal torque on the disc."
+ However. this is only valid up to the position in the dise where the orbits cross.," However, this is only valid up to the position in the disc where the orbits cross."
+ We adopt a torque function per unit disc mass on the disc of the form where g is a constant., We adopt a torque function per unit disc mass on the disc of the form where $g$ is a constant.
+ To mocel the tidal ellects of orbit crossings. we select the torque parameters so that the torque acts in the region where the particle orbits begin to cross.," To model the tidal effects of orbit crossings, we select the torque parameters so that the torque acts in the region where the particle orbits begin to cross."
+ We choose g=4 so that the disc is truncated. quickly., We choose $g=4$ so that the disc is truncated quickly.
+ The form of T4 controls where the dise is truncated., The form of $T_3$ controls where the disc is truncated.
+ We choose To model the effects of the inllowing cireumstellar gas. we model gas injection at some radius ring over a narrow region ol radial width 2i at a steady rate Alis with the local Ixeplerian speed.," We choose To model the effects of the inflowing circumstellar gas, we model gas injection at some radius $r_{\rm inj}$ over a narrow region of radial width $2 w$ at a steady rate $\dot M_{\rm inj}$ with the local Keplerian speed."
+Function S(r) describes themass injection that we take to be where £r) is unity for |r<] and zero otherwise.,Function $S(r)$ describes themass injection that we take to be where $H(x)$ is unity for $|x| < 1$ and zero otherwise.
+ We adopt a width i=0.0046(ri;DNE rH. ," We adopt a width $w = 0.0046 \,(r_{\rm inj}/r_{\rm H})^\frac{1}{2}\, r_{\rm H}$ "
+considerably faster (han ~O;:.,considerably faster than $\sim \Omega_K$.
+ In sum. we cannot present results on [frequenciesabove cQr)sco. and the form of our cooling function potentially suppresses some fluctuation power at the higher Ilrequency. end of the range we cdo discuss.," In sum, we cannot present results on frequenciesabove $\simeq 0.7 \nu_{ISCO}$, and the form of our cooling function potentially suppresses some fluctuation power at the higher frequency end of the range we do discuss."
+ Within the simulation. we do not consider any interaction between the emitted radiation and the material.," Within the simulation, we do not consider any interaction between the emitted radiation and the material."
+ ILowever. more realistically. there is always some opacity. ancl in most circumstances (he dominant opacity in (he material near a black hole is electron scattering.," However, more realistically, there is always some opacity, and in most circumstances the dominant opacity in the material near a black hole is electron scattering."
+ This opacity leads {ο a natural division of the radiation in (wo parts: that emitted inside or outside the photosphere., This opacity leads to a natural division of the radiation in two parts: that emitted inside or outside the photosphere.
+ Within the photosphere. scattering can add substantially to the lime a photon can take to reach the outside. washing oul fInctuations in intrinsic emissivitv: outside the photosphere. of course. scattering has very little effect on photon escape time.," Within the photosphere, scattering can add substantially to the time a photon can take to reach the outside, washing out fluctuations in intrinsic emissivity; outside the photosphere, of course, scattering has very little effect on photon escape time."
+" In addition. photons deriving their energy. [rom dissipation inside and outside (he photosphere can be distinguished spectrally: Inside (he photosphere. thermalization is strong. and the ocal spectrum should be approximately black body. at a temperature ~10 eV in AGN, ον] keV in GBIs."," In addition, photons deriving their energy from dissipation inside and outside the photosphere can be distinguished spectrally: Inside the photosphere, thermalization is strong, and the local spectrum should be approximately black body, at a temperature $\sim 10$ eV in AGN, $\sim 1$ keV in GBHs."
+ By contrast. outside (he photosphere. much lower gas densities and much veher ratios of heating density to mass density lead (o much higher temperatures. and the primary emission mechanism is inverse Compton scattering. so (hat the radiated spectrum is characteristically a power-law extending well into (he hard. X-ray regime.," By contrast, outside the photosphere, much lower gas densities and much higher ratios of heating density to mass density lead to much higher temperatures, and the primary emission mechanism is inverse Compton scattering, so that the radiated spectrum is characteristically a power-law extending well into the hard X-ray regime."
+ In order (o nake a realistic estimate of the light curve directly. from the simulations emissivity data. we therefore restrict our efforts to the coronal hard. X-ray emission. whose source is near or outside the scattering photosphere.," In order to make a realistic estimate of the light curve directly from the simulation's emissivity data, we therefore restrict our efforts to the coronal hard X-ray emission, whose source is near or outside the scattering photosphere."
+ To locate that photosphere requires a calculation of the opacity. vet its magnitude is not defined in code-units because the simulation requires no absolute density scale.," To locate that photosphere requires a calculation of the opacity, yet its magnitude is not defined in code-units because the simulation requires no absolute density scale."
+ Instead. we determine it alter the [act by the following procedure: We distinguish quantities in code-units from «quantiües in physical units by attaching a subscript e to the former. and leaving the latter unlabeled.," Instead, we determine it after the fact by the following procedure: We distinguish quantities in code-units from quantities in physical units by attaching a subscript $c$ to the former, and leaving the latter unlabeled."
+" [fa fraction 7 of the rest-mass of accretion were transformed into Ilumimosity ab infinity. i( would be because the unit of length is GA/c if G=e 1. and wá. when measured in units of ο, 15 dimensionless."," If a fraction $\eta$ of the rest-mass of accretion were transformed into luminosity at infinity, it would be because the unit of length is $GM/c^2$ if $G=c=1$ , and $u^\mu$, when measured in units of $c$, is dimensionless."
+ Here. y is the determinant of the metric.," Here, $g$ is the determinant of the metric."
+ Normalizing (he luminosity {ο the Eddington Iuminosity. Lj. we find that the relation between physical density.ancl code," Normalizing the luminosity to the Eddington luminosity $L_E$ , we find that the relation between physical densityand code"
+Soft CGamunatray Repeaters (SCRs) are a class of peculiar hiel-euergev sources cdiscovere through their recturent cluission of soft οταν bursts,Soft Gamma-ray Repeaters (SGRs) are a class of peculiar high-energy sources discovered through their recurrent emission of soft $\gamma$ -ray bursts.
+ These bursts have typical durations of Οι s and lhuuinosities iu the rauge 107-107 ores ! (sec Thuley2000 for a review of this class of objects)., These bursts have typical durations of $\sim$ 0.1 s and luminosities in the range $^{39}$ $^{42}$ ergs $^{-1}$ (see \cite{hurleyrew} for a review of this class of objects).
+" The bursting activity and the persistent emissiou observed iu the —0.5-10 keV enerev range are eenerallv explained in the framework of he ""Magnetar model (e.g. Duncan&Thompson 1992.. ""Thompson&Duncan 1995)). as caused bv ai highly naeguctized (D ~10 0) slowly rotating (P~ 5-8 x) ieutron star."," The bursting activity and the persistent emission observed in the $\sim$ 0.5-10 keV energy range are generally explained in the framework of the “Magnetar” model (e.g. \cite{dt92}, , \cite{td95}) ), as caused by a highly magnetized $B\sim$ $^{15}$ G) slowly rotating $P\sim$ 5-8 s) neutron star."
+ SCR 150620 is one of the most active Soft Gamunua- Repeaters., SGR 1806–20 is one of the most active Soft Gamma-ray Repeaters.
+ IElere we report uew observations of his source obtained with the INTEGRAL satellite in October 2003 during ai period of bursting activitv (Ciotzetal.2003a.. Ihulevetal. 2003.. Meoreeghettietal. 20035... Gotzetal. 20035).," Here we report new observations of this source obtained with the INTEGRAL satellite in October 2003 during a period of bursting activity \cite{gotza}, \cite{hurley}, \cite{merec}, \cite{gotzb}) )."
+ These data have two advantages compared to previous observations in the soft -cav onergv rauge of bursts from this source., These data have two advantages compared to previous observations in the soft $\gamma$ -ray energy range of bursts from this source.
+ First. they have been obtained with an miaenmg iustruineut. thus we can exclude that the bursts originate from a different source m the feld.," First, they have been obtained with an imaging instrument, thus we can exclude that the bursts originate from a different source in the field."
+ Second. thev have a eood seusitivitv and time resolution which allows us to study the spectral evolution of relatively ται! bursts.," Second, they have a good sensitivity and time resolution which allows us to study the spectral evolution of relatively faint bursts."
+ The region of was observed by INTEGRAL (Winkleretal. 20033). between October 8 and 15 2003 as part of the Core Program deep observation of he Galactic Centre. vielding an exposure of about [NO ks on the source.," The region of was observed by INTEGRAL \cite{winkler}) ) between October 8 and 15 2003 as part of the Core Program deep observation of the Galactic Centre, yielding an exposure of about $\sim$ 480 ks on the source."
+ Several bursts from the direction. of were detected i near real iue bv the INTECRAL Dust Alert System (IBÀS. Alereghettietal. 2003a)). using data from the IBIS instrmuent (Ubertinietal. 2003)).," Several bursts from the direction of were detected in near real time by the INTEGRAL Burst Alert System (IBAS, \cite{ibas}) ), using data from the IBIS instrument \cite{ibis}) )."
+ IBIS. a coded masis clescope with a large field of view (2907\297). colprises two detector lavers: ISGRI (15 keV. - 1 MeV. Lebrunetal. 20033) aud PICSIT (170 keV - 10. MeV. Labautietal. 2003)).," IBIS, a coded mask telescope with a large field of view $^{\circ}\times29^{\circ}$ ), comprises two detector layers: ISGRI (15 keV - 1 MeV, \cite{isgri}) ) and PICsIT (170 keV - 10 MeV, \cite{picsit}) )."
+ Ouly ISGRI data are relevaut here. since PICSIT docs not have enough time resolution for the study such short bursts.," Only ISGRI data are relevant here, since PICsIT does not have enough time resolution for the study such short bursts."
+ Iun total. 2] bursts were detected by the IBAS progranus.," In total, 21 bursts were detected by the IBAS programs."
+ By ueans of nuages accunmmlated over the time imtervals corresponding to the individual bursts. we can be confident that all of them originated from.," By means of images accumulated over the time intervals corresponding to the individual bursts, we can be confident that all of them originated from."
+. Iu fact. the derived coordinates are all within 2/ from the well known position of SCR 1506-20 (kaplanetal. 2002)). while the confidence error circle is typically of the order of 2.57.," In fact, the derived coordinates are all within $'$ from the well known position of SGR 1806-20 \cite{chandra}) ), while the confidence error circle is typically of the order of $'$."
+ Iu particular. the bursts positions are not consistent with the possible SCR. 1508-20. (Lambetal. 2003)) recently discovered at 15’ from.," In particular, the bursts positions are not consistent with the possible SGR 1808-20 \cite{lamb}) ) recently discovered at $'$ from."
+. The backeround subtracted light curves of the bursts. binned at 10 1s. are shown in Fig. 1..," The background subtracted light curves of the bursts, binned at 10 ms, are shown in Fig. \ref{lc}."
+ In order to merease the sigual-to-noise-ratio. they were extracted frou ISCRI pixels ihuninated by the source for at least half of their surface and selecting counts in the 15-100 keV energy range (nost of the bursts aad little or no signal at higher enerev).," In order to increase the signal-to-noise-ratio, they were extracted from ISGRI pixels illuminated by the source for at least half of their surface and selecting counts in the 15-100 keV energy range (most of the bursts had little or no signal at higher energy)."
+" The ""nursts were detected at various off-axis aueles. ranging from 2.5 o 123.5 deerees. corresponding ο à variation of in the instrument effective area;"," The bursts were detected at various off-axis angles, ranging from 2.5 to 13.3 degrees, corresponding to a variation of in the instrument effective area."
+ The ight curves shown in Fig., The light curves shown in Fig.
+ 1. have been corrected for this vienetting effect., \ref{lc} have been corrected for this vignetting effect.
+ The tota ΠΟΥ of net counts actually recorded for cach burstis iudicated in the correspondiug »uiel., The total number of net counts actually recorded for each burstis indicated in the corresponding panel.
+ The Πο curves shown in Fig., The light curves shown in Fig.
+ 1. have shapes typical or SCR bursts, \ref{lc} have shapes typical for SGR bursts.
+ From the light curves we determined the, From the light curves we determined the
+"most easily shown by ranking the planets in terms of their radius anomaly, i.e. the difference between its measured radius and the one predicted by models of a pure solar-composition planet of the same mass and age (?)..","most easily shown by ranking the planets in terms of their radius anomaly, i.e. the difference between its measured radius and the one predicted by models of a pure solar-composition planet of the same mass and age \citep{Guillot+06}."
+" As shown by Fig. 10,,"," As shown by Fig. \ref{fig:Ranomaly},"
+" when using standard models, CoRoT-2b has a positive, large radius anomaly of kkm, but still smaller than that of HAT-P-8b, TrES-4b, and WASP-12b."," when using standard models, CoRoT-2b has a positive, large radius anomaly of km, but still smaller than that of HAT-P-8b, TrES-4b, and WASP-12b."
+" However, for these last three planets, their large radius can be explained within the error bars by an additional heat source equivalent to of the incoming stellar luminosity deposited at the planet's center (see Guillot 2005 and references therein)."," However, for these last three planets, their large radius can be explained within the error bars by an additional heat source equivalent to of the incoming stellar luminosity deposited at the planet's center (see Guillot 2005 and references therein)."
+" As shown in the right panel of Fig. 10,,"," As shown in the right panel of Fig. \ref{fig:Ranomaly},"
+" this is not true for CoRoT-2b: because of its large mass, the planet tends to contract rapidly and therefore requires special conditions to explain its large size."," this is not true for CoRoT-2b: because of its large mass, the planet tends to contract rapidly and therefore requires special conditions to explain its large size."
+ We now consider how models of the planet's atmosphere affect its evolution., We now consider how models of the planet's atmosphere affect its evolution.
+" Remarkably, secondary transits of CoRoT-2b were detected in the optical from CoRoT lightcurves (??) and in the infrared from Spitzer IRAC observations (?) as well as ground-based measurements (?).."," Remarkably, secondary transits of CoRoT-2b were detected in the optical from CoRoT lightcurves \citep{Alonso+09, Snellen+10} and in the infrared from Spitzer IRAC observations \citep{Gillon+10} as well as ground-based measurements \citep{ADKR10}."
+ These directly probe the planetary atmosphere and are thus key constraints of the outer boundary conditions used in the evolution modeling., These directly probe the planetary atmosphere and are thus key constraints of the outer boundary conditions used in the evolution modeling.
+ The fluxes inferred from these measurements are equivalent to brightness temperaturesof 13252180 KK at 8 um and 1805+ KK at 4.5 um (??)..," The fluxes inferred from these measurements are equivalent to brightness temperaturesof $1325\pm 180$ K at $8\,\mu $ m and $1805\pm 70$ K at $4.5\,\mu $ m \citep{Gillon+10, Snellen+10}."
+ Additional ground-based measurements yield Ty=18901330 K in the Κ. band (2.1 um) and an upper limit of 2250K in the H band ( 1.6m) (?)..," Additional ground-based measurements yield $T_{\rm b}=1890^{+260}_{-350}\,$ K in the $_s$ band $\sim 2.1\,\mu$ m) and an upper limit of $2250\,$ K in the H band $\sim 1.6\,\mu$ m) \citep{ADKR10}. ."
+" In the optical, independent studies yield brightness temperatures that are very similar within error bars, i.e. 212079, K (22) and 2170+ 50K (?).."," In the optical, independent studies yield brightness temperatures that are very similar within error bars, i.e. $2120^{+90}_{-110}\,$ K \citep{Alonso+09,Alonso+10} and $2170\pm{50}\,$ K \citep{Snellen+10}."
+ A crucial consequence that can be derived is that The case of the optical brightness measurements are more complex to interpret because they may arise from either direct emission or a reflection of the incoming stellar flux (??)..," A crucial consequence that can be derived is that The case of the optical brightness measurements are more complex to interpret because they may arise from either direct emission or a reflection of the incoming stellar flux \citep{Alonso+09,Snellen+10}."
+" In the limit of a geometric albedo of 0.2, the flux would be entirely due to direct reflection and thus provide no information about the atmospheric temperature profile."," In the limit of a geometric albedo of $0.2$, the flux would be entirely due to direct reflection and thus provide no information about the atmospheric temperature profile."
+" In contrast, a zero albedo would imply that the flux in the optical is thermal emission from the planet, and that the corresponding temperatures are high."," In contrast, a zero albedo would imply that the flux in the optical is thermal emission from the planet, and that the corresponding temperatures are high."
+ These temperature constraints are compared in Fig., These temperature constraints are compared in Fig.
+ 11 to temperature profiles calculated for CoRoT-2b in the framework of our semi-gray analytical model assuming a full redistribution of the incoming stellar flux (seeEq. (6)) and9)., \ref{fig:atm-corot2b} to temperature profiles calculated for CoRoT-2b in the framework of our semi-gray analytical model assuming a full redistribution of the incoming stellar flux \citep[see Eq.~(\ref{eq:t4-global}) ) and.
+" The value of the intrinsic flux Tins= 1000K was chosen to match that of models with a size ~1.5 Rjyp, as observed."," The value of the intrinsic flux $\tint=1000\,$ K was chosen to match that of models with a size $\sim 1.5\,R_{\rm Jup}$ , as observed."
+ Usingvalues of the thermal and visible opacities calibrated, Usingvalues of the thermal and visible opacities calibrated
+The mass was determined separately for the 13.confirmed Centaurus group members. the M82 eroup members. (he (wo independent bins of galaxies (Din 1 and Bin 2). and for the larger number of probable satellites with each additional satellite galaxy. added in increasing radial projection from NGC 5128.,"The mass was determined separately for the 13 Centaurus group members, the M83 group members, the two independent bins of galaxies (Bin 1 and Bin 2), and for the larger number of probable satellites with each additional satellite galaxy added in increasing radial projection from NGC 5128."
+ The adopted value of 5 was determined independently for each list of galaxies in (he same manner as was done for the GCS in Section 3.1.. with tvpical values ranging from 5=1.5 to +=3.2.," The adopted value of $\gamma$ was determined independently for each list of galaxies in the same manner as was done for the GCS in Section \ref{mass:GCS}, with typical values ranging from $\gamma=1.5$ to $\gamma=3.2$."
+ The caleulated masses with projection corrections. are shown in Tables 2 3 and are plotted in Figure 7..," The calculated masses with projection corrections, are shown in Tables \ref{tab:IGG}
+ \ref{tab:GG} and are plotted in Figure \ref{fig:mass}."
+ The mass calculation for the eroup cannot be continued out indefinitely [ar because at some radius. the fundamental assumption of virial equilibrium breaks down.," The mass calculation for the group cannot be continued out indefinitely far because at some radius, the fundamental assumption of virial equilibrium breaks down."
+ At very. large radii. the group has not had enough time to undergo full relaxation (Cotéetal.1997).," At very large radii, the group has not had enough time to undergo full relaxation \citep{cote97}."
+. The ονπάσα aud crossing limes (Binnev&Tremaine1987).. when set equal to a IIubble time. occur al radii of 1967 kpe and 1400 kpc. respectively. for the entire G2 ealaxies surrounding NGC 5123.," The dynamical and crossing times \citep{binney87}, when set equal to a Hubble time, occur at radii of 1967 kpc and 1400 kpc, respectively, for the entire 62 galaxies surrounding NGC 5128."
+" Similarly. IXarachentsevοἱal.(2006) [ound a ""surface of zero velocity” around Centaurus of 1440 kpe. coinciding with the upturn in velocity dispersion. noted bv the cumulative galaxy bins seen in Fig."," Similarly, \cite{kar06} found a ”surface of zero velocity” around Centaurus of 1440 kpc, coinciding with the upturn in velocity dispersion noted by the cumulative galaxy bins seen in Fig."
+ 5 and discussed in Section 2.3.., \ref{fig:veldisp} and discussed in Section \ref{kin_disc}.
+ Thus an appropriate racial limit for (he mass caleulation is 2~1.5 Ape., Thus an appropriate radial limit for the mass calculation is $R\simeq1.5$ Mpc.
+ Thus. the large masses determined for the outer 15-20 galaxies. (includes Bin 2 and cumulative galaxy bins 42-62) are not valid as they are not likely virialized objects in the svstem.," Thus, the large masses determined for the outer 15-20 galaxies, (includes Bin 2 and cumulative galaxy bins 42-62) are not valid as they are not likely virialized objects in the system."
+ This is carried through to the mass-to-light ratios determined for these outer objects in Section 3.3.., This is carried through to the mass-to-light ratios determined for these outer objects in Section \ref{masstolight}. .
+" The mass-to-light ratios were calculated [rom the total mass. M,. determined in Section and a B-band luminosity and galactic extinction for each galaxy [rom (2004)."," The mass-to-light ratios were calculated from the total mass, $_t$, determined in Section \ref{mass} and a B-band luminosity and galactic extinction for each galaxy from \cite{kar04}."
+. All mass-to-light ratios are listed in Tables 2 3..," All mass-to-light ratios are listed in Tables \ref{tab:IGG}
+ \ref{tab:GG}."
+ The D-band luminosity for the ealaxv calculations include all of the galaxies within the outermost radii of each bin., The B-band luminosity for the galaxy calculations include all of the galaxies within the outermost radii of each bin.
+ The AI83 complex. therefore. contributes to the luminosity of Din 1 and Bin 2 in Table 2.. although it is not included as part of the mass tracer population.," The M83 complex, therefore, contributes to the luminosity of Bin 1 and Bin 2 in Table \ref{tab:IGG}, although it is not included as part of the mass tracer population."
+ The D-band Ipuninosity of NGC 5128 is 3.7950.01x10! L.., The B-band luminosity of NGC 5128 is $3.79\pm0.01 \times 10^{10}$ $_{\sun}$ .
+ This single galaxv makes up 66% of the entire light of the Centaurus eroup of galaxies., This single galaxy makes up $66\%$ of the entire light of the Centaurus group of galaxies.
+ The mass-to-light ratio of NGC 5128 alone is 5242? M. /L. out to R=45 kpe., The mass-to-light ratio of NGC 5128 alone is $52\pm22$ $_{\sun}$ $_{\sun}$ out to $R=45$ kpc.
+" The mass of the Centaurus groupof galaxies is determined trom this study to be ML, ont to the dynamical radius. leading to M/Ly,=153450 M. /L..These"," The mass of the Centaurus groupof galaxies is determined from this study to be $(9.2\pm3.0) \times 10^{12}$ $_{\sun}$ out to the dynamical radius, leading to $_B = 153\pm50$ $_{\sun}$$_{\sun}$ .These"
+Practically all star formation is thought to occur in clusters that are embedded in giant molecular clouds (GMCs) (2)..,Practically all star formation is thought to occur in clusters that are embedded in giant molecular clouds (GMCs) \citep{Ladas2003}.
+ This suggests that a clustering environment. where multiple objects compete for a common gus reservoir (22).. plays an important role in early stages of protostellar evolution. such as dictating the form of the stellar initial mass function (IMP) (2:: 2).," This suggests that a clustering environment, where multiple objects compete for a common gas reservoir \citep{Zinnecker1982, Larson1992}, plays an important role in early stages of protostellar evolution, such as dictating the form of the stellar initial mass function (IMF) \citealt{Bonnelletal1997}; ; \citealt{Bonnelletal2001b}) )."
+ Furthermore. ?. has collected observational evidence suggesting that star formation is a rapid orocess. Occurring on roughly the crossing time of the region at a variety of scales.," Furthermore, \citet{Elmegreen2000} has collected observational evidence suggesting that star formation is a rapid process, occurring on roughly the crossing time of the region at a variety of scales."
+ Not only does he propose that the star formation in atypical GMC occurs within 4 Myr (approximately the crossing ime for standard GMC) but that the cloud’s dispersal occurs within afew crossing times. or 10 Myr.," Not only does he propose that the star formation in a typical GMC occurs within 4 Myr (approximately the crossing time for standard GMC) but that the cloud's dispersal occurs within a few crossing times, or $\la$ 10 Myr."
+ The combined implication= of these observations is that star ormation occurs quickly and in groups and that the sites of star ormation disperse quickly., The combined implication of these observations is that star formation occurs quickly and in groups and that the sites of star formation disperse quickly.
+ Our proposal in this paper is that this is possible if GMCs are dynamically unbound objects. with the internal turbulent energy greater than that of the cloud’s gravity.," Our proposal in this paper is that this is possible if GMCs are dynamically unbound objects, with the internal turbulent energy greater than that of the cloud's self-gravity."
+ This follows from the work of ? who showed that transient (unbound) GMC sized objects can be formed from flows in the ISM., This follows from the work of \citet*{Semadenietal1995} who showed that transient (unbound) GMC sized objects can be formed from flows in the ISM.
+ We also find that unbound GMCs may provide a natural mechanism for the creation of OB associations. a notion first suggested by ?..," We also find that unbound GMCs may provide a natural mechanism for the creation of OB associations, a notion first suggested by \citet{Ambart1958}."
+ In the rest of this first section we discuss the ideas behind rapid star formation and the dynamical state of GMCs., In the rest of this first section we discuss the ideas behind rapid star formation and the dynamical state of GMCs.
+ We also include in this section a discussion of OB associations., We also include in this section a discussion of OB associations.
+ In section we describe the details of the simulation and section follows the general evolution of the GMC., In section we describe the details of the simulation and section follows the general evolution of the GMC.
+ In section+ we give estimates of the star formation efficiency in the GMC based on some simple assumptions., In section we give estimates of the star formation efficiency in the GMC based on some simple assumptions.
+ In section we highlight the similarities between the simulation and the general structure in an OB associations., In section we highlight the similarities between the simulation and the general structure in an OB associations.
+ A summary of the paper's main conclusions can be found in section6., A summary of the paper's main conclusions can be found in section.
+ Until the last decade or so GMCs were generally believed to be long-lived structures. with some estimates of ages reaching as high as 107 Myr (?22)..," Until the last decade or so GMCs were generally believed to be long-lived structures, with some estimates of ages reaching as high as $10^{8}$ Myr \citep*{Solomonetal1979, Scovilleetal1979, ScovilleHersh1979}."
+ Tt was generally believed that the chemistry of turning atomic species into molecules would require millions of years beforean object like a GMC would be detectable via its CO abundance (2).., It was generally believed that the chemistry of turning atomic species into molecules would require millions of years beforean object like a GMC would be detectable via its CO abundance \citep{Jura1975}.
+ One also had the problem that the CO mass in the galaxy. coupled with estimates of the star formation rate. suggested that GMCs had to live for tens of millions of years if the star formation efficiency was to remain at the observed level of a ew percent (22)..," One also had the problem that the CO mass in the galaxy, coupled with estimates of the star formation rate, suggested that GMCs had to live for tens of millions of years if the star formation efficiency was to remain at the observed level of a few percent \citep{ZuckermanEvans1974,
+ZuckermanPalmer1974}."
+ Recent observations of embedded clusters tend to suggest that he whole process of star formation. including GMC formation and dispersal. occurs on roughly the crossing time for the region (2).," Recent observations of embedded clusters tend to suggest that the whole process of star formation, including GMC formation and dispersal, occurs on roughly the crossing time for the region \citep{Elmegreen2000}."
+ ot only do most molecular clouds in the solar neighbourhood contain signs of star formation in the form of clusters. but the age determination of these clusters suggests they are very young. ypically lessthan IOMvr (2)...," Not only do most molecular clouds in the solar neighbourhood contain signs of star formation in the form of clusters, but the age determination of these clusters suggests they are very young, typically lessthan 10Myr \citep{Hartmann2000}. ."
+ This suggests that star formation, This suggests that star formation
+"spectroscopy, surface brightness fluctuations, or resolved stellar populations.","spectroscopy, surface brightness fluctuations, or resolved stellar populations."
+" In order to not complicate the analysis with large stellar mass-to-light ratio variations, we also include onlyold objects (ages Z; 5 where possible, e.g., excluding some young, extended Gyr)clusters that are known in the LMC and beyond."," In order to not complicate the analysis with large stellar mass-to-light ratio variations, we also include only objects (ages $\ga$ 5 Gyr) where possible, e.g., excluding some young, extended clusters that are known in the LMC and beyond."
+" The results are plotted in Figure 8,, with the details, references, and full data table provided in Appendix ??.."," The results are plotted in Figure \ref{fig:uber}, with the details, references, and full data table provided in Appendix \ref{sec:uber}."
+ There are various interesting features in the size-luminosity parameter space., There are various interesting features in the size-luminosity parameter space.
+" The most compact objects include the classical population of GCs with ry 2-4 pc, the UCDs which extend up to ry~ 50 pc, and the fainter ECs up to ry 25 pc, which comprise not only extended clusters (Huxoretal.2005,2008,2009,2011a;Mackey 2011),, but also faint fuzzies (Larsen&Brodie2000;Larsen2002;Santosetal.2007;Scheepmaker *diffuse star clusters"" (Pengetal.2006),, and the 2007),,Palomar clusters in the Milky Way halo."," The most compact objects include the classical population of GCs with $r_{\rm h}\sim$ 2–4 pc, the UCDs which extend up to $r_{\rm h}\sim$ 50 pc, and the fainter ECs up to $r_{\rm h}\sim$ 25 pc, which comprise not only extended clusters \citep{2005MNRAS.360.1007H,2008MNRAS.385.1989H,2009ApJ...698L..77H,2011MNRAS.414..770H,2006ApJ...653L.105M,2008AJ....135.1482S,2010MNRAS.404.1157M,2011ApJ...730..112C,2011ApJ...738...58H}, but also faint fuzzies \citep{2000AJ....120.2938L,2002AJ....124.1410B,2005A&A...442...85S,2006ApJ...638L..79H,2008AJ....135.1567H,2007A&A...467.1003C,2007A&A...469..925S}, “diffuse star clusters"" \citep{2006ApJ...639..838P}, and the Palomar clusters in the Milky Way halo."
+" There is then an apparent gap between galaxies and star clusters, which is now seen to be a diagonal region rather than a simple size gap (cf the “Shapley line"" of vandenBergh2008;; and also Gilmoreetal.2007; 2011))."," There is then an apparent gap between galaxies and star clusters, which is now seen to be a diagonal region rather than a simple size gap (cf the “Shapley line” of \citealt{2008MNRAS.390L..51V}; and also \citealt{2007ApJ...663..948G,2008MNRAS.389.1924F,2009ARA&A..47..371T,2011MNRAS.414.3699M}) )."
+" This gap corresponds roughly to a line of constant surface brightness, ryοςL-V?, so there may be a selection effect at work here, with deeper imaging and spectroscopic surveys needed."," This gap corresponds roughly to a line of constant surface brightness, $r_{\rm h} \propto L^{-1/2}$, so there may be a selection effect at work here, with deeper imaging and spectroscopic surveys needed."
+" There are a few, possibly rare, “bridging” objects between the UCDs and the cEs, which as we discuss below, may imply that star clusters and galaxies are not completely distinct populations."," There are a few, possibly rare, “bridging” objects between the UCDs and the cEs, which as we discuss below, may imply that star clusters and galaxies are not completely distinct populations."
+" Considering the UCDs and noting the typical strong selection biases against objects fainter than My~—11.5, we see that the data suggest a nearly flat size trend that parallels the compact GCs."," Considering the UCDs and noting the typical strong selection biases against objects fainter than $M_V \sim -11.5$, we see that the data suggest a nearly flat size trend that parallels the compact GCs."
+" The previous paradigm of a strong size-luminosity trend for UCDs is trumped by the new discoveries of low-luminosity UCDs, mostly from around M87 but also from a few other galaxies."," The previous paradigm of a strong size-luminosity trend for UCDs is trumped by the new discoveries of low-luminosity UCDs, mostly from around M87 but also from a few other galaxies."
+" This includes the Milky Way where the halo cluster NGC 2419 was long thought to be a unique object, while it can now be seen as a harbinger of the new class of UCDs (see black square in the lower panel Figure 8))."," This includes the Milky Way where the halo cluster NGC 2419 was long thought to be a unique object, while it can now be seen as a harbinger of the new class of UCDs (see black square in the lower panel Figure \ref{fig:uber}) )."
+" We thus see that standard identifications of UCDs by luminosity alone are inadvisable, as they coexist with compact GCs over a factor of ~ 15 range in luminosity."," We thus see that standard identifications of UCDs by luminosity alone are inadvisable, as they coexist with compact GCs over a factor of $\sim$ 15 range in luminosity."
+" The best that can be done in the absence of direct size information is to estimate the probability of a given object being a UCD or a GC (cf Figure 3)), with objects more luminous than My~—12.5 fairly safely designated as UCDs."," The best that can be done in the absence of direct size information is to estimate the probability of a given object being a UCD or a GC (cf Figure \ref{fig:hist}) ), with objects more luminous than $M_V \sim -12.5$ fairly safely designated as UCDs."
+" Interestingly, both giant GCsand stripped nuclei are expected to follow strong size-luminosity relations and yet we find no such relation for the M87 (r,—L)UCDs."," Interestingly, both giant GCs stripped nuclei are expected to follow strong size-luminosity $L$ ) relations and yet we find no such relation for the M87 UCDs."
+" Upon closer scrutiny over the full range of luminosities, the Τμ], relation for nuclei is not particularly strong (see Section ??))."," Upon closer scrutiny over the full range of luminosities, the $L$ relation for nuclei is not particularly strong (see Section \ref{sec:cmd2}) )."
+" The lack of an r,—L relation for UCDs", The lack of an $L$ relation for UCDs
+The epoch of reionization in hydrogen has become a topic of considerable interest (Barkana Loch 2001: Fan. Cavill. Keating 2006: Meiksin 2009: Furlanetto 22009) as a probe of the transition from neutral to ionized hwdrosen iu the intergalactic medium (ICAL).,"The epoch of reionization in hydrogen has become a topic of considerable interest (Barkana Loeb 2001; Fan, Carilli, Keating 2006; Meiksin 2009; Furlanetto 2009) as a probe of the transition from neutral to ionized hydrogen in the intergalactic medium (IGM)."
+" This transition occurred somewhere between redshifts 2=612. marking the exit from the cosmic ""dark ages”. beeinning at the time when the first stars and galaxies formed at redshifts 5>30 (Teemark 1997: Ricotti, Cunuediu. Shull 2002: Treuti Shull 2010)."," This transition occurred somewhere between redshifts $z =$ 6–12, marking the exit from the cosmic “dark ages"", beginning at the time when the first stars and galaxies formed at redshifts $z > 30$ (Tegmark 1997; Ricotti, Gnedin, Shull 2002; Trenti Shull 2010)."
+ Helium underwent sinular reionization from tto (Chat is. from Ue! to IHe!2) at :=2:80.2 (Reimers 11997: Shull 22001). imiost. likely mediated by the harder (GE>51.4 eV) radiation frou quasars and other active galactic nuclei (ACN).," Helium underwent similar reionization from to (that is, from $^{+}$ to $^{+2}$ ) at $z = 2.8 \pm 0.2$ (Reimers 1997; Shull 2004), most likely mediated by the harder $E \geq 54.4$ eV) radiation from quasars and other active galactic nuclei (AGN)."
+ With a {ανα ionization potential. Te! is harder to ionize than IP. aud Ie!? recombines 5G times faster than II! (Osterbrock Ferland 2006: Fardal. Girous. Shull 1998).," With a 4 ryd ionization potential, $^+$ is harder to ionize than $^0$, and $^{+2}$ recombines 5--6 times faster than $^+$ (Osterbrock Ferland 2006; Fardal, Giroux, Shull 1998)."
+ For these reasous. aud the fact that most hot stars lack strong [| rvd contiuua. it is believed that ACN are the primary agents of rreionization.," For these reasons, and the fact that most hot stars lack strong 4 ryd continua, it is believed that AGN are the primary agents of reionization."
+" Touization models find that. owing to its resilience. Ile! is meh more abundant than IT"". with predicted colunim-deusity ratios 4j= 0) zz KO100 (NMiralda-Escudé 11996: Fardal. Cüroux. Shull 1995)."," Ionization models find that, owing to its resilience, $^+$ is much more abundant than $^0$, with predicted column-density ratios $\eta \equiv$ ) $\approx$ 50–100 (Miralda-Escudé 1996; Fardal, Giroux, Shull 1998)."
+ Determining when aud how the universe was ionized has been an important question in cosmology for decades (Camm Peterson 1965: Sunvaev 1977)., Determining when and how the universe was ionized has been an important question in cosmology for decades (Gunn Peterson 1965; Sunyaev 1977).
+ Although recent progress has led to quantitative coustraiuts. we still do not kuow whether galaxies are the sole agents of lydrogen reionization. aud the epoch of reiomizatiou remains uncertain.," Although recent progress has led to quantitative constraints, we still do not know whether galaxies are the sole agents of hydrogen reionization, and the epoch of reionization remains uncertain."
+ Wedo know that lydrogen reionization in the IGM is complete by zz6. based on the rapid evolution between redshifts +=6.2 and 2ο of CCumu-Petersou (GP) absorption from neutral hydrogeu along lines of sight to QSOs (Fan 22002: Cuedin 2001).," We know that hydrogen reionization in the IGM is complete by $z \approx 6$, based on the rapid evolution between redshifts $z = 6.2$ and $z \approx 5$ of Gunn-Peterson (GP) absorption from neutral hydrogen along lines of sight to QSOs (Fan 2002; Gnedin 2004)."
+" The optical depth of the cosmic mucrowave backerouud (I&omatsu 22010) sets aso coustraiut ou instantaneous reiouization at 2,>OS aud a lo confidence mterval of 2,=LO541.2.", The optical depth of the cosmic microwave background (Komatsu 2010) sets a $3 \sigma$ constraint on instantaneous reionization at $z_r > 6.5$ and a $1 \sigma$ confidence interval of $z_r = 10.5 \pm 1.2$.
+" A linut 2,27 is set by the detection of celuitters at +=6.5 Guajor evolution iu the observed luminositv function would be expected if neutral eas", A limit $z_r \geq 7$ is set by the detection of emitters at $z = 6.5$ (major evolution in the observed luminosity function would be expected if neutral gas
+time reveal the high proper motion of tle object.,time reveal the high proper motion of the object.
+ A small but resolved uuebula is also detected., A small but resolved nebula is also detected.
+ Optical inaging has been obtained ou two occasions with the Subaru telescope in 1999 and 2003. aud iu 2001 by the UST.," Optical imaging has been obtained on two occasions with the Subaru telescope in 1999 and 2003, and in 2001 by the HST."
+ Data from the first epoch have been retrieved from the SMOIKA database at the eud of January 2003., Data from the first epoch have been retrieved from the SMOKA database at the end of January 2003.
+ These first observations were obtained between April 22 and May 15 1999 1sing the SuprimcCam istrict (Alivazalki et al. 19958 )), These first observations were obtained between April 22 and May 15 1999 using the SuprimeCam instrument (Miyazaki et al. \cite{miyazaki1998}) )
+ mounted at the Casseerain focus., mounted at the Cassegrain focus.
+ This configuration wuch used a 3x2 CCD setup was in operation until Juv 1999 to perform an initial detailed testing of both tl telescope aud iustruneut., This configuration which used a 3x2 CCD setup was in operation until July 1999 to perform an initial detailed testing of both the telescope and instrument.
+ Each ST-002A SITe CCD hac 2018 x 1096 square pixels with a size of 15g corres20ving to oon the «kv at the Casscerain focus., Each ST-002A SITe CCD had 2048 x 4096 square pixels with a size of $\mu$ corresponding to on the sky at the Cassegrain focus.
+ For our purpose. we ouly used CCD 5 which fully coutains the field of3219.," For our purpose, we only used CCD 5 which fully contains the field of."
+ Of the 3 filters used. D. R aud broad. there exist only fia-ficlds for tI BR baud.," Of the 3 filters used, B, R and broad, there exist only flat-fields for the R band."
+ The ffter L65] has a central wavelcneth of aadxl a full wiclti of325À., The filter $\_$ L651 has a central wavelength of and a full width of.
+ Due to relatively large central waveleisth variations with off axis anele (up to 100 AJ. this fiter was later decommissioned.," Due to relatively large central wavelength variations with off axis angle (up to 100 ), this filter was later decommissioned."
+ Laboratory lueasurenmients show tiat the làιο was in al cases well insice the filter baud pass but die to the spaial ial1noseneities the exact transmission at rremiains uncerzin., Laboratory measurements show that the line was in all cases well inside the filter band pass but due to the spatial inhomogeneities the exact transmission at remains uncertain.
+ Tadividual exposure times were ss for the aadxl D filters and ss for t1ο KR band., Individual exposure times were s for the and B filters and s for the R band.
+ The telescope was nuoved by abou iin both directions οποσα two consecutive ex]x191110, The telescope was moved by about in both directions between two consecutive exposures.
+", A] individual images were corrected for bias usii8 over-scan areas and the BR baud nuages were furthermore fat-fielded using well illuuinated dome flats.", All individual images were corrected for bias using over-scan areas and the R band images were furthermore flat-fielded using well illuminated dome flats.
+ Since only one bias exposure exiss for these observations. we prefered to correct imuages σας mica over-scan values in order to not increase pliotcmetric axd yositional errors.," Since only one bias exposure exists for these observations, we prefered to correct images using mean over-scan values in order to not increase photometric and positional errors."
+ Individual nuages were the1 moved to a colon frame using a well exposed sinee reference star located close to the neutron star positkπι in order O nuumuiize the effects due to the lack of correction or geometrical distortion., Individual images were then moved to a common frame using a well exposed single reference star located close to the neutron star position in order to minimize the effects due to the lack of correction for geometrical distortion.
+ The images were fjen stacked using a statistical cosuic-rav event rejection criterion axd finally rebiuned with a ppixel size., The images were then stacked using a statistical cosmic-ray event rejection criterion and finally rebinned with a pixel size.
+ The IIST archive provided secoud epoch οservations., The HST archive provided second epoch observations.
+ We used the ταiftered CCD (5ICCD aperture) mage obtained on July 2] 2001 which offers the highest sensitivity., We used the unfiltered CCD (50CCD aperture) image obtained on July 21 2001 which offers the highest sensitivity.
+ Πιοvidual TST inages were ecometrically corrected. drizz outo a reference maüaese with a pixel scale half of t1ο original one. filtered. for cosmic-ray Παρατς aud sacked together following the method outlined iu FruchY Took (2 021).," Individual HST images were geometrically corrected, drizzled onto a reference image with a pixel scale half of the original one, filtered for cosmic-ray impacts and stacked together following the method outlined in Fruchter Hook \cite{fruchter2002}) )."
+ These data were preseuted in Ίναau et al. (2003a))., These data were presented in Kaplan et al. \cite{kaplan03a}) ).
+ Third epoch nuages were all obtained on June 8 2003 with the FOCAS iustruimnent (Ixashikwwa oet al. 20121} , Third epoch images were all obtained on June 8 2003 with the FOCAS instrument (Kashikawa et al. \cite{kashikawa2002}) )
+operated at the CasseerainC» focus., operated at the Cassegrain focus.
+ The targetOo was located on CCD 2 owuch has the best cosmetic quality., The target was located on CCD 2 which has the best cosmetic quality.
+ We obtaired series of D aud R baud exposures wihn inteeration times of )ss and 6003 respectively., We obtained series of B and R band exposures with integration times of s and s respectively.
+ Raw nuages were corrected fex bias using average offset exposures and flafielded wihi dome flats., Raw images were corrected for bias using average offset exposures and flat-fielded with dome flats.
+ A dedicated IDL procedure then COLTEEed. individual images for ecolmetrical distorion which oheqwise could lead to errors ofup to aat the ee ofthe field of view., A dedicated IDL procedure then corrected individual images for geometrical distortion which otherwise could lead to errors of up to at the edge of the field of view.
+ All images were loved to a COMMON yale usiie a set of 9 reference stcllar- objects and tji stacked. together using a statistical oulerion for rojectiusg cosnic-ray 1npacts., All images were moved to a common frame using a set of 9 reference stellar-like objects and then stacked together using a statistical criterion for rejecting cosmic-ray impacts.
+ Cromund based οservations were in eeneral otained in photometric coxditious althougi sone thin cimus nav have been preseut πι 2003., Ground based observations were in general obtained in photometric conditions although some thin cirrus may have been present in 2003.
+ The Subaru telescope is equipped with an atnospleric refraction correctioL priui at the telescope side., The Subaru telescope is equipped with an atmospheric refraction correction prism at the telescope side.
+ The differeice du colours between the relatively red astrometric reference stars aud the blue uettrou star thus docs not vield aiv significant difference 1n position., The difference in colours between the relatively red astrometric reference stars and the blue neutron star thus does not yield any significant difference in position.
+ The log of observatious is listed in Table 1.., The log of observations is listed in Table \ref{obslog}.
+ FWIIM secings are given as measured on the suuunied images., FWHM seeings are given as measured on the summed images.
+ Diving the second SulXU €observations. D aud R photometric standard stars were repeatedly observed iu the field of SALTIQ0-501. (Liuolt. 1992)).," During the second Subaru observations, B and R photometric standard stars were repeatedly observed in the field of SA110-504 (Landolt \cite{landolt}) )."
+ Ilowever. zero point variability. by about L1 mag rnis reveals the probable preseuce of thin cirrus during that üght.," However, zero point variability by about 0.1 mag rms reveals the probable presence of thin cirrus during that night."
+ We thus used our photometric CEITT images obtained in 1998 axd described im Motch et al. L999 ), We thus used our photometric CFHT images obtained in 1998 and described in Motch et al. \cite{motch1999}) )
+ to calibrate several Ba xd R local photometric reference stars in adiition to those already metioned in that paoer., to calibrate several B and R local photometric reference stars in addition to those already mentioned in that paper.
+ Conipariic CFUT aud Subaru 2003 zero poiuts shows that the 2003 images were indeed. absbed by ~ 0.15 mag in D aud BR. MaguiticCR wore COPted using a 2-d Gaussian fitting oxocess taki iuto account possible slaw backerouud spatial variabili, Comparing CFHT and Subaru 2003 zero points shows that the 2003 images were indeed absorbed by $\sim$ 0.15 mag in B and R. Magnitudes were computed using a 2-d Gaussian fitting process taking into account possible sky background spatial variability.
+ The maenuiudes of 3219. listed in Table 2 are eeneralv consistent witi those reported by Kaplan et al. (20032) j," The magnitudes of , listed in Table \ref{magrx} are generally consistent with those reported by Kaplan et al. \cite{kaplan03a}) )"
+ from IIST/STIS photometry., from HST/STIS photometry.
+ Errors quoted are one siena aic take mto accotnt uucertainties iu colour ranstormaion aud for the 1999 D value errors due to the absence of flat-Beld. correction., Errors quoted are one sigma and take into account uncertainties in colour transformation and for the 1999 B value errors due to the absence of flat-field correction.
+ There is uo evidence for Hux variability oj finie scaC8 ο [vears, There is no evidence for flux variability on time scales of years.
+ We show in Fie., We show in Fig.
+ 1l the sununed D nuage obtained in 2003 at the Subaru clescope with FOCAS., \ref{sub03} the summed B image obtained in 2003 at the Subaru telescope with FOCAS.
+ Using the bes determined value in 2003 we find R 10.32 4 0.17., Using the best determined value in 2003 we find $-$ R = +0.32 $\pm$ 0.17.
+ This colour iudex is significantly redder han that of R= 0.61 £0.13. van Ierkwijk ντκανα 20 Ha) aud maybe also R= O35 +OL EKuluni van Ierkwijk L998)}).," This colour index is significantly redder than that of $-$ R = $-$ 0.61 $\pm$ 0.13, van Kerkwijk Kulkarni \cite{vk2001}) ) and maybe also $-$ R = $-$ 0.3 $\pm$0.4, Kulkarni van Kerkwijk \cite{KvK98}) )."
+derived primarily [rom proper motion surveys. iuay be deficient in youug. low space-1uotion cwarfs.,"derived primarily from proper motion surveys, may be deficient in young, low space-motion dwarfs."
+ Figure 12. also compares our analyses with the star formation liistorles proposed [rou C dwar studies by Barry(1988) aud Rocha-Pintoetal.(2000)., Figure \ref{fig-sfhist} also compares our analyses with the star formation histories proposed from G dwarf studies by \citet{barry} and \citet{rp00}.
+. The poor time resolutiou at large age:. coupled with the lack of calibrating clusters and tie relatively sinall size of the VC sample. limiS the utility of comparisous at older ages ( >3 Civis).," The poor time resolution at large ages, coupled with the lack of calibrating clusters and the relatively small size of the VC sample, limits the utility of comparisons at older ages ( $>3$ Gyrs)."
+ Nonetheless. it is encouraging that all of tle moclels predict similar [ractions of stars younger tlan d Cyrs.," Nonetheless, it is encouraging that all of the models predict similar fractions of stars younger than 4 Gyrs."
+ At vounger ages. the M dwarf data are better calibrated.," At younger ages, the M dwarf data are better calibrated."
+ There is uo evideuce [rom otw analysis for the substantia uumbers of young G dwarls ascribed by Barry(1988). to a recent birst of star formation., There is no evidence from our analysis for the substantial numbers of young G dwarfs ascribed by \citet{barry} to a recent burst of star formation.
+ Simila‘ly. our data fail to match the details of the Rocha-Pintoetal.(2000) analysis.," Similarly, our data fail to match the details of the \citet{rp00} analysis."
+" ludeed. Rocha-Piutoοἱal.(2000) s ""Burst A7 of voung (350.5 Gyr) GC dwarfs is accotipaniecd by a delicieucy of M. cdwarls. while their ""AB Gap” at 1-2 Gyrs corresponds to an apparent excess of M. dwarfs. (altlough. as we noted above. we believe this feature reflects a delicieucy in our analysis metLoc raljer than a burst of star formation)."," Indeed, \citet{rp00}' 's “Burst A” of young $\lesssim 0.5$ Gyr) G dwarfs is accompanied by a deficiency of M dwarfs, while their “AB Gap” at 1-2 Gyrs corresponds to an apparent excess of M dwarfs (although, as we noted above, we believe this feature reflects a deficiency in our analysis method rather than a burst of star formation)."
+" The ""AB"" eap found by Rocha-Pintoetal.(2000) correspouds to le Vaughau-Prestou gap Vaughan&P‘estou{1950).", The “AB” gap found by \citet{rp00} corresponds to the Vaughan-Preston gap \citet{vp80}.
+. As noted above (Section L.1)). Herbst&Mile(1989) |ave suggested hat the sparse 1tunber of early ype M dw:urls with weak enission (oir Groip C) imielt be an analogous feaure.," As noted above (Section \ref{activity}) ), \citet{hm89} have suggested that the sparse number of early type M dwarfs with weak emission (our Group C) might be an analogous feature."
+ However. observatious of Hyades M dwarls show that the Πε| limit correspouds o TiO»5-0.55 in hat cluster.," However, observations of Hyades M dwarfs show that the $\alpha$ limit corresponds to $\sim$ 0.55 in that cluster."
+" Tlis. the weak eimmission M εἶννα[s in Grou »Cnui st have ages of less han the age «M tie Hyacles (0.6 €Wis} — Le. iges while1 matcl Rocha-Piitoetal. (2000)s ""Burst A."," Thus, the weak emission M dwarfs in Group C must have ages of less than the age of the Hyades (0.6 Gyrs) – i.e. ages which match \citet{rp00}' 's “Burst A”."
+ As Soderloi1.Duncan.&Joinsou(L991) ancl Rex:lia-Piitoetal.(2000) rave discussed. oue ieecis either a COLiplicated G-dwarl age-activivy relatio 110 Ina cha coustant sta‘formation history. Or a complicaed sta ‘formation listory to save the simple C-«warf age-activity relation.," As \citet{sdj91} and \citet{rp00} have discussed, one needs either a complicated G-dwarf age-activity relation to match a constant star formation history, or a complicated star formation history to save the simple G-dwarf age-activity relation."
+ Given the ack of agree.101 |etween our À cdwarl auaVses alk the CC dwarf analyses. we believe tliat he complicated €iedwar. age-activity relation exdlanatiol is [avored.," Given the lack of agreement between our M dwarf analyses and the G dwarf analyses, we believe that the complicated G-dwarf age-activity relation explanation is favored."
+ The existence of a large spread lu activity 1n he coeval C clwarls of M67 (CH€ampapaetal.2000) aud the large spread iu roalion ‘ates in stars of youig clusters (Barnes1907 sIggesMODs that auy age-activity relation Is coiiplex. with stochastic star O slar variations.," The existence of a large spread in activity in the coeval G dwarfs of M67 \citep{giampapa_m67} and the large spread in rotation rates in stars of young clusters \citep{barnes} suggests that any age-activity relation is complex, with stochastic star to star variations."
+ Iludeec. AL dwars show similar beliaviour. with a suitering of dMe dwarfs bluer than the Ho limi| dn σοιje clusters.," Indeed, M dwarfs show similar behaviour, with a smattering of dMe dwarfs bluer than the $\alpha$ limit in some clusters."
+ Biuarity may well be a contributiug acOr at both spectral types., Binarity may well be a contributing factor at both spectral types.
+ Other measures of the recent 5ar fonvation history have been made without reference to chromospheric activity., Other measures of the recent star formation history have been made without reference to chromospheric activity.
+ Hernandez.Valls-Gabaud.&Gilmore(2000) have used Hipparcos maguituce diagrams to derive he St ‘formation liistory of the Solar Neighibourlood within the last 3 Gyrs., \citet{x00} have used Hipparcos color-magnitude diagrams to derive the star formation history of the Solar Neighbourhood within the last 3 Gyrs.
+ They find au oscilatory «Οι;»oneut of star formation with a period of ().5 Gyr superposed on an underlyiug coustau star foruation rate., They find an oscillatory component of star formation with a period of 0.5 Gyr superposed on an underlying constant star formation rate.
+ Their Figure [| inclicates that tley see roughly 2.5 times as much star formation f'om 1.6—2.6 Cyr as at ~0.3—1.0 Cyr., Their Figure 4 indicates that they see roughly 2.5 times as much star formation from $\sim 1.6-2.6$ Gyr as at $\sim 0.3 - 1.0$ Gyr.
+ It is suggestive that ταῖς is similar to the deviations f‘Ol conSant star formation seen in our Figure 12 - hat is. the possible delicieucy of young stars.," It is suggestive that this is similar to the deviations from constant star formation seen in our Figure \ref{fig-sfhist} – that is, the possible deficiency of young stars."
+ Nettier our data nor Hernandez.Valls-Gabaud.&Gilmore(2000) show the lull iu star formation yetweel. ||-2 Gyrs seen by Rocha-Pintoetal. (2000).., Neither our data nor \citet{x00} show the lull in star formation between 1-2 Gyrs seen by \citet{rp00}. .
+likely be small ancl all-sky. sensitive. arcminute resolution experiments will be needed for their detection.,"likely be small and all-sky, sensitive, arcminute resolution experiments will be needed for their detection."
+ Ilere the spherical wavelets will be tested against non-Gaussian simulations of artificially specified moments that will be assumed to be small., Here the spherical wavelets will be tested against non-Gaussian simulations of artificially specified moments that will be assumed to be small.
+ In this case a useful way to construct non-Gaussian distributions is by perturbing the Gaussian one throughDo a sum of moments. the Edgeworth5 expansion.," In this case a useful way to construct non-Gaussian distributions is by perturbing the Gaussian one through a sum of moments, the Edgeworth expansion."
+ For simplicity we will consider the two lowest cumulants to characterise the deviations from normality: skewness and kurtosis., For simplicity we will consider the two lowest cumulants to characterise the deviations from normality: skewness and kurtosis.
+ As discussed in the introduction alternative models to stancarc inflation. e.g. cosmic defects as a subdominant source of density. perturbations or non-stancard inflation. can produce significant levels of at least one of the two moments.," As discussed in the introduction alternative models to standard inflation, e.g. cosmic defects as a subdominant source of density perturbations or non-standard inflation, can produce significant levels of at least one of the two moments."
+ For small deviations from Ciaussianitv. there is a wide class of distributions that can be given in terms of a Gaussian distribution times an infinite sum of its cumulants.," For small deviations from Gaussianity, there is a wide class of distributions that can be given in terms of a Gaussian distribution times an infinite sum of its cumulants."
+ This is the well known Edgeworth expansion., This is the well known Edgeworth expansion.
+ The problem with this expansion is that setting all cumulants to zero except one does not guarantee the positive definiteness and normalization that a distribution has to satisfy., The problem with this expansion is that setting all cumulants to zero except one does not guarantee the positive definiteness and normalization that a distribution has to satisfy.
+ However. for small deviations from normality the resulting function is always positive at least up to many sigmas in the tail of the distribution and the normalization factor required for the function to become a well defined cistribution is very small and does not appreciably disturb the non-zero moments (i.e. skewness or kurtosis) introduced in the first place.," However, for small deviations from normality the resulting function is always positive at least up to many sigmas in the tail of the distribution and the normalization factor required for the function to become a well defined distribution is very small and does not appreciably disturb the non-zero moments (i.e. skewness or kurtosis) introduced in the first place."
+" The Edgeworth expansion can be obtained. [rom the characteristic Function ó(/) by considering the linear terms in the cumulants and. performing the inverse Fourier transform to recover the density function. f(r): where df, is the Llermite polvnomial.", The Edgeworth expansion can be obtained from the characteristic function $\phi(t)$ by considering the linear terms in the cumulants and performing the inverse Fourier transform to recover the density function $f(x)$: where $H_n$ is the Hermite polynomial.
+ Considering the perturbations corresponding to the skewness ancl kurtosis and keeping only the first terms in the corresponding Lermite polynomials. we have where S. Ix denote skewness ancl kurtosis. respectively.," Considering the perturbations corresponding to the skewness and kurtosis and keeping only the first terms in the corresponding Hermite polynomials, we have where S, K denote skewness and kurtosis, respectively."
+ We will use these equations to generate our artificially specified non-Gaussian distributions., We will use these equations to generate our artificially specified non-Gaussian distributions.
+ Since the resulting distribution is not well defined even. for the case of small skewness and kurtosis we set the function to zero when it becomes negative and we also normalize it appropriately., Since the resulting distribution is not well defined even for the case of small skewness and kurtosis we set the function to zero when it becomes negative and we also normalize it appropriately.
+ We remark that the zero cuts of the distribution. if present. appear far away in the tails of the distribution for the case of small values of skewness and kurtosis that we consider here.," We remark that the zero cuts of the distribution, if present, appear far away in the tails of the distribution for the case of small values of skewness and kurtosis that we consider here."
+ Also. as à consequence. the normalization value required is very ‘lose to 1.," Also, as a consequence, the normalization value required is very close to 1."
+ In this wav we checked that the initial values of je skewness and kurtosis we start with in the Ecdeeworth *xpansion does not appreciably change after the necessary hanges introduced. to obtain a well defined: probability ensity function (pef)., In this way we checked that the initial values of the skewness and kurtosis we start with in the Edgeworth expansion does not appreciably change after the necessary changes introduced to obtain a well defined probability density function (pdf).
+ 1n order to make the simulations resemble the CMD ata observed by a given experiment we smooth them with a Gaussian filter., In order to make the simulations resemble the CMB data observed by a given experiment we smooth them with a Gaussian filter.
+ For practical reasons we use a FWIHIAL of 33. which may correspond to some ol the channels in all-sky experiments like NLAP and Planck (c.g. the θα Lz channe of the Planck mission)., For practical reasons we use a FWHM of $33'$ which may correspond to some of the channels in all-sky experiments like MAP and Planck (e.g. the $30$ GHz channel of the Planck mission).
+ We choose to work on the LUIEALPix pixelisation of the sphere with a resolution Nya=256., We choose to work on the HEALPix pixelisation of the sphere with a resolution $N_{side}=256$.
+ We use the HEALDPix package to perform the analysis of our simulated CMD cata., We use the HEALPix package to perform the analysis of our simulated CMB data.
+ However. it is not. adequate to use that package to convolve our unfiltered. independen temperature data with the Gaussian 33° PWLAL beam in Fouricr space. instead we perform the convolution in rea space.," However, it is not adequate to use that package to convolve our unfiltered independent temperature data with the Gaussian $33'$ FWHM beam in Fourier space, instead we perform the convolution in real space."
+ After that. in order to make the simulations more realistic we normalize the CMD power spectrum C' of both Gaussian and non-Gaussian simulations to that of a CDA Hat A-mocdel using the HIZALDix package.," After that, in order to make the simulations more realistic we normalize the CMB power spectrum $C_l$ of both Gaussian and non-Gaussian simulations to that of a CDM flat $\Lambda$ -model using the HEALPix package."
+ As a consequence of the beam convolution and the introduction of correlations in the temperature maps the original levels of skewness ancl kurtosis injected through the Edgeworth expansion are reduced. (compare columns 1 and 2 in table 2))., As a consequence of the beam convolution and the introduction of correlations in the temperature maps the original levels of skewness and kurtosis injected through the Edgeworth expansion are reduced (compare columns 1 and 2 in table \ref{Fisher}) ).
+ The performance of spherical wavelets will be tested with these simulations in section 5., The performance of spherical wavelets will be tested with these simulations in section 5.
+ Since wavelet coefficients. represent. linear. transformations of the original data. in the case of a Gaussian distribution the wavelet cocllicicnts remain Gaussian distributed.," Since wavelet coefficients represent linear transformations of the original data, in the case of a Gaussian distribution the wavelet coefficients remain Gaussian distributed."
+ This a very nice property of wavelets and all we have to do to test CGaussianity in wavelet space is to look from deviations from normality., This a very nice property of wavelets and all we have to do to test Gaussianity in wavelet space is to look from deviations from normality.
+ l]lowever. for the case of the sphere any given pixelisation scheme will introduce biases.," However, for the case of the sphere any given pixelisation scheme will introduce biases."
+ The specific bias introduced. will depend on. for instance. whether the pixels are not of equal area or the distances between one pixel and its neighbours vary with the position on the sphere.," The specific bias introduced will depend on, for instance, whether the pixels are not of equal area or the distances between one pixel and its neighbours vary with the position on the sphere."
+ This is in fact the situation for the two pixelisations already used (ο analyse all-skv οΑΠΣ temperature Uuetuations., This is in fact the situation for the two pixelisations already used to analyse all-sky CMB temperature fluctuations.
+ For the COBLE-DAIR: experiment the pixelisation usec was the Quac-Cube and in this projection of the cube on the sphere equal-area. pixels on the sides of the cube appear with different: area when projected on the sphere., For the COBE-DMR experiment the pixelisation used was the Quad-Cube and in this projection of the cube on the sphere equal-area pixels on the sides of the cube appear with different area when projected on the sphere.
+ For present satcllite experiments like SLAP ancl Planck the LIEALPix pixclisation is now widely used., For present satellite experiments like MAP and Planck the HEALPix pixelisation is now widely used.
+ While this. pixelisation possesses very nice properties. such as equal area iso-atitude pixels. however the distances between one pixel and. its neighbours vary with latitude.," While this pixelisation possesses very nice properties, such as equal area iso-latitude pixels, however the distances between one pixel and its neighbours vary with latitude."
+ Pixels near the equator tend to be more uniformly distributed than those near the poles., Pixels near the equator tend to be more uniformly distributed than those near the poles.
+ As we will compute in next section. this property. produces a bias in the kurtosis of the wavelet coellicients for the case of the SIIW. (see table 1. Gaussian case which corresponds to a null injected. value for the kurtosis).," As we will compute in next section, this property produces a bias in the kurtosis of the wavelet coefficients for the case of the SHW (see table 1, Gaussian case which corresponds to a null injected value for the kurtosis)."
+ For the Gaussian, For the Gaussian
+interaction sites between SNRs aud molecular clouds.,interaction sites between SNRs and molecular clouds.
+" By correlating SNR. uasers with known GeV aud Τον sources, we have identified an emerge class of 5-raxv- bright interacting SNRs."," By correlating SNR masers with known GeV and TeV sources, we have identified an emerging class of $\gamma$ -ray-bright interacting SNRs."
+ Of the 21 known Maser SNRs currently identified there are ten with Ce to TeV-ΟΠΟΙΟΥ οταν associations. and six with both.," Of the 24 known Maser SNRs currently identified there are ten with GeV to TeV-energy $\gamma$ -ray associations, and six with both."
+ If +-rav cCluission from these sources is larecly duc to hadronic COSMIC raves. the enhanced local cosmic rav lPonization rates in these clouds can explain the production of OII iuolecules behind a C-type shock. suggested by Wardle(1999).," If $\gamma$ -ray emission from these sources is largely due to hadronic cosmic rays, the enhanced local cosmic ray ionization rates in these clouds can explain the production of OH molecules behind a C-type shock, suggested by \citet{wardle99}."
+. Furthermore. cosmüc rav loulzation is typically comparable to or donünant over ionizations from iuterior thermal A-ravs. though without detailed knowledge of the cosimic ray spectzüm these results have large uncertainties.," Furthermore, cosmic ray ionization is typically comparable to or dominant over ionizations from interior thermal X-rays, though without detailed knowledge of the cosmic ray spectrum these results have large uncertainties."
+ Interacting SNRs represent a pronusing class of 5-rav sources which are likely to be uncovered by future 5-rav observatories., Interacting SNRs represent a promising class of $\gamma$ -ray sources which are likely to be uncovered by future $\gamma$ -ray observatories.
+provided only short. widely spaced runs lor sampling nova rates in external galaxies.,"provided only short, widely spaced runs for sampling nova rates in external galaxies."
+ We have attempted to overcome this bias bv using a dedicated telescope to observe the target galaxy in its entirety. every clear night for many months.," We have attempted to overcome this bias by using a dedicated telescope to observe the target galaxy in its entirety, every clear night for many months."
+ Our first survey of this ivpe was of M81 (Neill&Shara2004). and produced a bulk nova rate higher than previous studies (Shafter.Ciardullo.&Pritehet2000)., Our first survey of this type was of M81 \citep{nei04} and produced a bulk nova rate higher than previous studies \citep{sha00}.
+. ILowever. we also demonstrated the effects of dust in the disk on the detection of novae in the bulge. implving that our higher bulk rate could be still be low by up to a factor of 2.," However, we also demonstrated the effects of dust in the disk on the detection of novae in the bulge, implying that our higher bulk rate could be still be low by up to a factor of 2."
+ Nearby dwarl ellipticals offer relatively dust-DIree targets that could potentially be surveved for novae with close to completeness., Nearby dwarf ellipticals offer relatively dust-free targets that could potentially be surveyed for novae with close to completeness.
+ Yet. if one examines the plot of normalized nova rate versus galaxy. lunminositv. such as presented in the references above. it is clear (hat at the low luminosity end. there is much uncertainty.," Yet, if one examines the plot of normalized nova rate versus galaxy luminosity, such as presented in the references above, it is clear that at the low luminosity end, there is much uncertainty."
+ This is for obvious reasons., This is for obvious reasons.
+ In particular. low Iuminositv svstems produce few novae per vear and so the sample is small.," In particular, low luminosity systems produce few novae per year and so the sample is small."
+ The low Iuminositw svstenis miust be nearby aud so are often. very. large and clilficult to survey in their entirely (e.g. M33. LAIC. SAIC).," The low luminosity systems must be nearby and so are often very large and difficult to survey in their entirety (e.g. M33, LMC, SMC)."
+ We (ook advantage of the availability of telescope time on an hourly basis. provided by the Tenagra observatory. to perform a comprehensive. niehtly survey of four local group dwarl galaxies with the aim of refining the nova rates al the low luminosity end.," We took advantage of the availability of telescope time on an hourly basis, provided by the Tenagra observatory, to perform a comprehensive, nightly survey of four local group dwarf galaxies with the aim of refining the nova rates at the low luminosity end."
+ We surveved \I32. NGC! 205. NGC I47. and NGC 185 for over flour months every clear night.," We surveyed M32, NGC 205, NGC 147, and NGC 185 for over four months every clear night."
+ We are also surveving the LMC with a dillerent telescope. and will present the results from that survey im a subsequent paper.," We are also surveying the LMC with a different telescope, and will present the results from that survey in a subsequent paper."
+ These surveys will continue for several vears and will provide accurate nova rates Lor the low Iuminosity svstems. allowing us to determine if there is indeed a trend in nova rate with luminosity.," These surveys will continue for several years and will provide accurate nova rates for the low luminosity systems, allowing us to determine if there is indeed a trend in nova rate with luminosity."
+ In order to constrain binary formation and evolution theory. (his kind of survev must be accompanied by comprehensive. densely lime-sampled survevs of hieher mass galaxies.," In order to constrain binary formation and evolution theory, this kind of survey must be accompanied by comprehensive, densely time-sampled surveys of higher mass galaxies."
+ Onlv by removing svstematie biases can we determine if there is a universal, Only by removing systematic biases can we determine if there is a universal
+distribution of subsamples of IRDCs in both the Ist and 4th quadrant of the Galactic plane (Simon et al.,distribution of subsamples of IRDCs in both the 1st and 4th quadrant of the Galactic plane (Simon et al.
+ 2006; Jackson et al., 2006; Jackson et al.
+ 2008; Marshall et al., 2008; Marshall et al.
+ 2009)., 2009).
+" Both kinematic and dust extinction techniques have been used to infer these distances, and although they lead to similar results, there are some differences (Fig. 4))."," Both kinematic and dust extinction techniques have been used to infer these distances, and although they lead to similar results, there are some differences (Fig. \ref{dist}) )."
+ The properties of these distance distributions are summarised in Table 1.., The properties of these distance distributions are summarised in Table \ref{tab:dist}.
+" Using dust extinction, Marshall et al. ("," Using dust extinction, Marshall et al. ("
+"2009) found a centrally peaked, Gaussian-like distance distribution very similar for both the 1st and Ath quadrant of the Galaxy (Fig.","2009) found a centrally peaked, Gaussian-like distance distribution very similar for both the 1st and 4th quadrant of the Galaxy (Fig."
+ 4 - bottom panel)., \ref{dist} - bottom panel).
+" In a similar way, kinematic shown in Fig."," In a similar way, kinematic shown in Fig."
+ 4 (top panel) show a good agreement between the 1st and 4th quadrant., \ref{dist} (top panel) show a good agreement between the 1st and 4th quadrant.
+ Although in the 1st quadrant tail at 5 kpc clearly emerges., Although in the 1st quadrant a tail at 5 kpc clearly emerges.
+" However most significanta difference between the extinction and kinematic distances is the position of the peak, being located at ~3 kpc in one case and at ~5 kpc in the other."," However most significant difference between the extinction and kinematic distances is the position of the peak, being located at $\sim3$ kpc in one case and at $\sim5$ kpc in the other."
+" Both techniques have their own biases∙ and advantages, it ∙∙is therefore difficult∙ to favor one distance distribution over another."," Both techniques have their own biases and advantages, it is therefore difficult to favor one distance distribution over another."
+ However a Gaussian distribution is a rather good approximation to the distance distribution in both quadrants., However a Gaussian distribution is a rather good approximation to the distance distribution in both quadrants.
+ Fig., Fig.
+ 4 (right) shows the average distance to the IRDCs as a function of galactic longitude for the sources with distances measured by these two techniques., \ref{dist} (right) shows the average distance to the IRDCs as a function of galactic longitude for the sources with distances measured by these two techniques.
+ Any systematic trend with longitude could introduce a bias in analysis adopting a statistical distribution for the distance of IRDCs., Any systematic trend with longitude could introduce a bias in analysis adopting a statistical distribution for the distance of IRDCs.
+ It is clear that there is very little variation of the IRDC distances with respect to the galactic longitude., It is clear that there is very little variation of the IRDC distances with respect to the galactic longitude.
+" The only region where there may be such an effect is towards galactic centre seems, an area which is not covered in our IRDC Spitzer catalogue which only extends into |=+10°."," The only region where there may be such an effect is towards galactic centre seems, an area which is not covered in our IRDC Spitzer catalogue which only extends into $l=\pm10\degr$."
+" It is worth noting that distance variations around the mean as a function of longitude are very similar for both methods, emphasising that is it predominantly only the average distance which differs between the two methods."," It is worth noting that distance variations around the mean as a function of longitude are very similar for both methods, emphasising that is it predominantly only the average distance which differs between the two methods."
+ Another possible bias is with respect to the size of the IRDCs., Another possible bias is with respect to the size of the IRDCs.
+ The studies shown in Fig., The studies shown in Fig.
+ 4 do not contain IRDCs as small as those in the Spitzer based sample and so it is possible that the small and large clouds have different distance distributions., \ref{dist} do not contain IRDCs as small as those in the Spitzer based sample and so it is possible that the small and large clouds have different distance distributions.
+" However,? recent observations with ATNF Mopra in CS J—1-0 (Peretto et al."," However, recent observations with ATNF Mopra in CS J=1-0 (Peretto et al."
+ in prep) of a square degree of the galactic plane, in prep) of a square degree of the galactic plane
+a neutron star while the short [lived systems (with Fa<100 Mars) did: undergo common envelope episodes. with hypereritical aceretion onto the neutron star.,a neutron star while the short lived systems (with $t_{\rm grav}< 100$ Myrs) did undergo common envelope episodes with hypercritical accretion onto the neutron star.
+ This common envelope episodes had two consequences: they tightened the orbits ancl Led to decrease of the mass ratio of the final double neutron star system. as one of the neutron stars. accreted some matter.," This common envelope episodes had two consequences: they tightened the orbits and led to decrease of the mass ratio of the final double neutron star system, as one of the neutron stars accreted some matter."
+ This implied that the mass distribution of the eravitational wave selected: sample of double neutron star binaries was cdillerent than the radio selected one., This implied that the mass distribution of the gravitational wave selected sample of double neutron star binaries was different than the radio selected one.
+ Llere we present a comparison of the radio selected sample of double neutron star binaries. taking into account selection ellects. with the gravitational wave selected. one.," Here we present a comparison of the radio selected sample of double neutron star binaries, taking into account selection effects, with the gravitational wave selected one."
+ The radio population contains also the binaries with merger imes above the Llubble time., The radio population contains also the binaries with merger times above the Hubble time.
+ In Figure 5. we present the distributions of expected objects in the plane spanned bv the mass ratio anc the oinmury mass defined as the more massive component. of a binary for APDO5. HIP and SP. models.," In Figure \ref{fig:expmass} we present the distributions of expected objects in the plane spanned by the mass ratio and the primary mass defined as the more massive component of a binary for APD05, HP and SP models."
+ The left. panels correspond to the radio selected: sample while the right vanels show the gravitational wave selected. ones., The left panels correspond to the radio selected sample while the right panels show the gravitational wave selected ones.
+ The objects contained in the radio sample are weighted by he time they are observable as pulsar from Earth. while he objects in the gravitational wave selected: sample are weighted by the volume in which they are observable.," The objects contained in the radio sample are weighted by the time they are observable as pulsar from Earth, while the objects in the gravitational wave selected sample are weighted by the volume in which they are observable."
+ The solid lines correspond. to constant values of the chirp mass in these coordinates., The solid lines correspond to constant values of the chirp mass in these coordinates.
+ Observed binary neutron stars are shown as black dots., Observed binary neutron stars are shown as black dots.
+ In addition on Figure 6 we show the chirp mass distributions of the binaries selected by their observability in the radio band. and by. their observability in gravitational waves.," In addition on Figure \ref{fig:ratio} we show the chirp mass distributions of the binaries selected by their observability in the radio band, and by their observability in gravitational waves."
+ Comparing our radio selected sample with the data rom Table Lowe see that the APDOS (top panel) and SP (lower panel) models are more consistent. with observations han the LIP (middle. panel) model in which majority of παν radio pulsars are predicted to have q«O.S aux one massive neutron star with ni1.6A4.., Comparing our radio selected sample with the data from Table 1 we see that the APD05 (top panel) and SP (lower panel) models are more consistent with observations than the HP (middle panel) model in which majority of binary radio pulsars are predicted to have $q< 0.8$ and one massive neutron star with $m_2 > 1.6 M_{\odot}$.
+ The APDOS model. which reproduces best the observed. distribution of oulsars in the 7?D diagram. is consistent with most of he observed binary radio pulsars.," The APD05 model, which reproduces best the observed distribution of pulsars in the $P-\dot P$ diagram, is consistent with most of the observed binary radio pulsars."
+ Lt does also reproduce very well the chirp mass distribution of the radio observec sample., It does also reproduce very well the chirp mass distribution of the radio observed sample.
+ Phe binary neutron stars with low mass ratios ane moderate masses are not predicted in this model., The binary neutron stars with low mass ratios and moderate masses are not predicted in this model.
+ However. one must take into account the fact that measurements of masses of JISII-1736 and J1518]4904 carry quite large error. bars.," However, one must take into account the fact that measurements of masses of J1811-1736 and J1518+4904 carry quite large error bars."
+ Ehe. SP. model. reprocuces the. distribution of masses and mass ratios of all observed. neutron. star binaries., The SP model reproduces the distribution of masses and mass ratios of all observed neutron star binaries.
+ In this model the distribution of masses is wicle., In this model the distribution of masses is wide.
+ Lt predicts existence of binaries containing massive neutron stars. however the racdo selected distribution is concentrated around qOS0.9 and. primary masses 1.35M...," It predicts existence of binaries containing massive neutron stars, however the radio selected distribution is concentrated around $q\sim 0.8-0.9$ and primary masses $\sim 1.35 M_{\odot}$."
+ A comparison of radio selected. sample with the eravitational wave one shows visible dillerences for SP and II models and negligible for the APDOS model (sec also Figure 6))., A comparison of radio selected sample with the gravitational wave one shows visible differences for SP and HP models and negligible for the APD05 model (see also Figure \ref{fig:ratio}) ).
+ In the model APDOS. the two samples are very similar as expected.," In the model APD05, the two samples are very similar as expected."
+ This is due to the fact that in the binary evolution model A the range of masses of newborn neutron stars is narrow comparing to the previous calculations as well the amount of matter acercted during the common envelope episode with a helium star is negligib, This is due to the fact that in the binary evolution model A the range of masses of newborn neutron stars is narrow comparing to the previous calculations as well the amount of matter accreted during the common envelope episode with a helium star is negligible.
+ In model LLP the radio sample is dominated by unequal mass binaries with q=0.7. mozzLl.7M..," In model HP the radio sample is dominated by unequal mass binaries with $q=0.7$, $m_2\approx 1.7
+M_\odot$."
+ Llowever the eravitational wave selected sample the equal mass binaries with both components of about 1.447. are also dominant., However the gravitational wave selected sample the equal mass binaries with both components of about $1.4 M_{\odot}$ are also dominant.
+ Also the chirp mass distributions are very different. with radio sample leaning towards the higher chirp mass binaries.," Also the chirp mass distributions are very different, with radio sample leaning towards the higher chirp mass binaries."
+ This result which is apparently counterintuitive. follows from the [aet that in this model we allow for substantial accretion and consequently strong reeveling of the first born pulsar.," This result which is apparently counterintuitive, follows from the fact that in this model we allow for substantial accretion and consequently strong recycling of the first born pulsar."
+ This pulsar has low magnetic field ancl remains radio loud for a very long time. and thus it contributes significantly to the radio sample.," This pulsar has low magnetic field and remains radio loud for a very long time, and thus it contributes significantly to the radio sample."
+ At the same time the binary has a large chirp mass because of the strong accretion., At the same time the binary has a large chirp mass because of the strong accretion.
+ The gravitational wave selected. sample is. dominated. in model HP by the short lived systems with no significant accretion episodes. and therefore contains more svstems with a low chirp mass.," The gravitational wave selected sample is dominated in model HP by the short lived systems with no significant accretion episodes, and therefore contains more systems with a low chirp mass."
+ Adeitionally this model fails to reproduce the radio observed. distribution of masses and mass ratios., Additionally this model fails to reproduce the radio observed distribution of masses and mass ratios.
+ In the SP model the gravitational wave selected sample is shifted toward the higher chirp mass values with respect to the radio selected one., In the SP model the gravitational wave selected sample is shifted toward the higher chirp mass values with respect to the radio selected one.
+" This is because of the volume effect. due to the Lact that the sampling volume scales as x,MET."," This is because of the volume effect, due to the fact that the sampling volume scales as $\propto {\cal M}^{5/2}$."
+ Thus the heavier binarics are observable from. much larger volume in gravitational waves., Thus the heavier binaries are observable from much larger volume in gravitational waves.
+ The mass ratio distribution of the gravitational wave selected sample leans toward lower values because the unequal mass binaries typically contain a low mass neutron star with a more massive companion., The mass ratio distribution of the gravitational wave selected sample leans toward lower values because the unequal mass binaries typically contain a low mass neutron star with a more massive companion.
+ Thus they have a higher chirp mass than the equal mass neutron star binaries that twpically contain two stars with low masses., Thus they have a higher chirp mass than the equal mass neutron star binaries that typically contain two stars with low masses.
+ This is similar to the comparison between the gravitational wave selected and. the volume. limited sample presented in Bulik et al., This is similar to the comparison between the gravitational wave selected and the volume limited sample presented in Bulik et al.
+ 2004., 2004.
+ Phe comparison with radio observations reveals quite and interesting feature: the distribution of chirp mass is underestimated in. this model. however the observed mass ratios are relatively well reproduced. as this model Leads to a wide distribution of racio observed. mass ratios.," The comparison with radio observations reveals quite and interesting feature: the distribution of chirp mass is underestimated in this model, however the observed mass ratios are relatively well reproduced, as this model leads to a wide distribution of radio observed mass ratios."
+ The small dilferences between the gravitational. wave selected. and the radio selected. distributions of mass ratio and chirp mass in model APDOS are mainly due to the fact that the initial masses of neutron stars come from a very narrow range and we do not allow for substantial accretion., The small differences between the gravitational wave selected and the radio selected distributions of mass ratio and chirp mass in model APD05 are mainly due to the fact that the initial masses of neutron stars come from a very narrow range and we do not allow for substantial accretion.
+ Thus most neutron star binaries have similar masses aud there is no possibility to form binaries with cdillerent. chirp masses or mass ratios in this mocel., Thus most neutron star binaries have similar masses and there is no possibility to form binaries with different chirp masses or mass ratios in this model.
+ Lowever. once we allow for a wide distribution of neutron star initial masses or for à significant accretion the dilferences between the two samples start to appear.," However, once we allow for a wide distribution of neutron star initial masses or for a significant accretion the differences between the two samples start to appear."
+spin is likely to be high (?)..,spin is likely to be high \citep{Heger:03sn}.
+ Therefore. since the kick velocity increases with ducreasing spin. this shall lead to a conservative survival probability for cach iierger tree.," Therefore, since the kick velocity increases with increasing spin, this shall lead to a conservative survival probability for each merger tree."
+ The survival probability shall be studied as a function of the initial mass of the seed IMDIT., The survival probability shall be studied as a function of the initial mass of the seed IMBH.
+ We choose a range of 13000AL.. to euconipass the plausible IMDBII formation clhamnels aud relevant masses involved., We choose a range of $10-3000 M_\odot$ to encompass the plausible IMBH formation channels and relevant masses involved.
+" For the three formation channels discussed in οον, the ποσα masses are likely to be the following: ~LOOOAL.. for a single stellar runaway (7): ~LOOAS.. for the stellar mass DII collision channel. where the seed IMDIT is produced by a dmassive supernova remnant or a small stellar runaway (2): and ~200LOOAL.. for Pop III remnants within a deuse globulur cluster (??).."," For the three formation channels discussed in \ref{sec:intro}, the seed masses are likely to be the following: $\sim 1000
+M_\odot$ for a single stellar runaway \citep{Portegies:2004fm}; $\sim
+100 M_\odot$ for the stellar mass BH collision channel, where the seed IMBH is produced by a massive supernova remnant or a small stellar runaway \citep{Heger:03sn}; and $\sim 200-400 M_\odot$ for Pop III remnants within a dense globular cluster \citep{Wise:05snpop3,Madau:01Pop3}."
+ One of the biggest uncertainties is determining a proper distribution for the secondary BIT masses., One of the biggest uncertainties is determining a proper distribution for the secondary BH masses.
+ As fgue 2 shows. the retention probability depends strouglv on the mass ratio between the IMDIT and DIL," As figure \ref{Fig-ret-prob} shows, the retention probability depends strongly on the mass ratio between the IMBH and BH."
+ Theoretical black hole mass distributions from solar unctallicity field populations tend to peak strougly around 1037... (7).., Theoretical black hole mass distributions from solar metallicity field populations tend to peak strongly around $10 M_\odot$ \citep{2001ApJ...554..548F}.
+ There are. however. several strong conrpetine effects in a. primordial elobular cluster that can change this distribution (e.g. low metallicity stellar evolution. high binary fraction. stellar collisions. lass scerceation. and natal kicks).," There are, however, several strong competing effects in a primordial globular cluster that can change this distribution (e.g., low metallicity stellar evolution, high binary fraction, stellar collisions, mass segregation, and natal kicks)."
+ For our fiducial experiment. we assunie that the secondary masses are selected from à Ixoupa IME with anu upper mass cutoff of 12037... (2)..," For our fiducial experiment, we assume that the secondary masses are selected from a Kroupa IMF with an upper mass cutoff of $120 M_\odot$ \citep{Kroupa:01imf}."
+ Recall that this vields an average stellar iuass of about LAL... much smaller than the Sspected average stellar iuass from a zero-nmietallicitv environment.," Recall that this yields an average stellar mass of about $1 M_\odot$, much smaller than the suspected average stellar mass from a zero-metallicity environment."
+ We further assume that each star above LOAL.. evolves directly to a DII with no mass loss., We further assume that each star above $10 M_\odot$ evolves directly to a BH with no mass loss.
+ This siuplified treatment givese us an average BIT mass of ~ 20A..., This simplified treatment gives us an average BH mass of $\sim 20 M_\odot$ .
+ Clearly. iiass loss would decrease the average DII nass even in these low metallicity priniordial elobulars.," Clearly, mass loss would decrease the average BH mass even in these low metallicity primordial globulars."
+" Therefore. we used a more sophisticated black hole mass ""uctiou that approximates the results of model Cl of Ux"," Therefore, we used a more sophisticated black hole mass function that approximates the results of model C1 of \citet{Belczynski:2006bh}."
+ This model includes the effect of mass loss from superuovae and winds for a population of stars with uctallicity Z=0.0001. though it has fewer massive remnants from binary mergers (cf Figure 8 of 2)) than is expected for a primordial globular cluster.," This model includes the effect of mass loss from supernovae and winds for a population of stars with metallicity $Z = 0.0001$, though it has fewer massive remnants from binary mergers (cf Figure 8 of \citet{Belczynski:2006bh}) ) than is expected for a primordial globular cluster."
+ Figure 3. clemonstrates the effect on the kick velocity distribution between these two black hole mass functions or a LOOOAL.. INIBIT., Figure \ref{fig:vkickbhmf} demonstrates the effect on the kick velocity distribution between these two black hole mass functions for a $1000 M_\odot$ IMBH.
+ Since the distribution of secondary basses is so uncertain. we demonstrated the effect of varvius the mass ratio in Figure 2..," Since the distribution of secondary masses is so uncertain, we demonstrated the effect of varying the mass ratio in Figure \ref{Fig-ret-prob}."
+ For a ncear-solar uctallicity stellar cluster. the entire population of black roles nav. be less than AZS20M (?)..," For a near-solar metallicity stellar cluster, the entire population of black holes may be less than $M \ltsim 20~M_\odot$ \citep{2001ApJ...554..548F}."
+" Finally, we mist select the eccentricity distribution."," Finally, we must select the eccentricity distribution."
+ If these INMDBITI-DII imersers were 2-body processes. we uieht expect the eccentricity of the orbit to be very jiearlv circular right before the black holes meree as eravitational radiation eriuds away the orbital angular Ποιοτα (?)..," If these IMBH-BH mergers were 2-body processes, we might expect the eccentricity of the orbit to be very nearly circular right before the black holes merge as gravitational radiation grinds away the orbital angular momentum \citep{Peters:1963ux}."
+ However. few body eucouuters are mich nore conumnuon Within a elobulu cluster because the interaction cross-section is much larger (?)..," However, few body encounters are much more common within a globular cluster because the interaction cross-section is much larger \citep{Heggie:book}."
+ Therefore. uauv Ες are shepherded iuto mergers with an IMDIT hrough exchange with lower mass black holes (?).. and he resulting eccentricity can be quite large (?)..," Therefore, many BHs are shepherded into mergers with an IMBH through exchange with lower mass black holes \citep{Miller:2002pi}, and the resulting eccentricity can be quite large \citep{Gultekin:2006tb}."
+ Iu fact. simulations have shown that rare interactions can vield nergers with e>0.999. aud such a highly ecceutric orbit can become even eccentric through. eravitational radiation emission (77).," In fact, simulations have shown that rare interactions can yield mergers with $e>0.999$, and such a highly eccentric orbit can become even eccentric through gravitational radiation emission \citep{Peters:1963ux,Kennefick:1998ab}."
+ Therefore. though rare.," Therefore, though rare,."
+"systems, To assign ecceutricities to cach merecr. we use the simulation results of ὃν, which clupirically characterizes the ecceutricity distribution as a function of the mass ratio of the encounter."," To assign eccentricities to each merger, we use the simulation results of \citet{Gultekin:2006tb}, which empirically characterizes the eccentricity distribution as a function of the mass ratio of the encounter."
+ Note. though. that equation (1)) is really ouly valid iu the μια]. eccentricity regiue and. thus. it may be true that the kicks can be even higher than those studied here for such nearly racial orbits.," Note, though, that equation \ref{eqn:Fit}) ) is really only valid in the small eccentricity regime and, thus, it may be true that the kicks can be even higher than those studied here for such nearly radial orbits."
+ Tn order to determine the probability that au IMDBII survives the short merger epoch after formation. our shumlatious cousist of 109 Monte Carlo realizatious of an N-step merger chain.," In order to determine the probability that an IMBH survives the short merger epoch after formation, our simulations consist of $10^6$ Monte Carlo realizations of an N-step merger chain."
+ The merger chains are tailored to munuc the initial conditions aud eucouuters predicted by current INIBIT formation theories within proto-eglobular and stellar cluster euvironnents as described in the previous section., The merger chains are tailored to mimic the initial conditions and encounters predicted by current IMBH formation theories within proto-globular and stellar cluster environments as described in the previous section.
+ During a ierecr. eravitational waves radiate not only lear momentum. but also angular momentum ancl enerev or mass.," During a merger, gravitational waves radiate not only linear momentum, but also angular momentum and energy or mass."
+ Fully relativistic nuuerical simulations sugecst that ~25% of the augular iunomentum (defined at the 3uneriuost stable circular orbit) can be radiated οπήιο the merger (?777)..," Fully relativistic numerical simulations suggest that $\sim 25\%$ of the angular momentum (defined at the innermost stable circular orbit) can be radiated during the merger \citep{Pretorius:2005gq,Campanelli05a,Baker:2006yw,Herrmann:2006ks}."
+ Du addition. the mass of the merecr product is ouly ~95% the mass of the two progenitor black holes (2???)..," In addition, the mass of the merger product is only $\sim 95\%$ the mass of the two progenitor black holes \citep{Pretorius:2005gq,Campanelli05a,Baker:2006yw,Herrmann:2006ks}."
+ ence. after cach step within a particular merecr tree. we adjust the total spin and imass of the merger remnant to account for these losses.," Hence, after each step within a particular merger tree, we adjust the total spin and mass of the merger remnant to account for these losses."
+ Figure 5 demonstrates the reteutiou perceutage after 25 mergers with black holes selected from à Ixoupa IMP as a function of the initial INIT mass., Figure \ref{fig:retain} demonstrates the retention percentage after 25 mergers with black holes selected from a Kroupa IMF as a function of the initial IMBH mass.
+ This figure indicates that retaining au IMDII of less than 10007. occurs less than 33% of the tine for the given distribution of black hole masses., This figure indicates that retaining an IMBH of less than $1000 M_\odot$ occurs less than $33\%$ of the time for the given distribution of black hole masses.
+ While the large number of the lower-nass black holes dominate the total umber of merecrs. he rare mergers with massive stellar-anass black holes dominate the ejections — and this treud only strengtheus as the INMDIT mass increases.," While the large number of the lower-mass black holes dominate the total number of mergers, the rare mergers with massive stellar-mass black holes dominate the ejections – and this trend only strengthens as the IMBH mass increases."
+ For example. figure { shows hat for a 1000.AL. seed. nearly all ejections come from lack holes with mass AL>30AL. and that most come from those with mass AJ~τοAL...," For example, figure \ref{fig:ejectfrac} shows that for a $1000~\Msun$ seed, nearly all ejections come from black holes with mass $M > 30~\Msun$ and that most come from those with mass $M \sim 70~\Msun$."
+ Naturally. his implies that if 230AL.. black holes are extremely rare in primordial globular clusters. the retention fraction Increases dramatically (7).. making gravitational recoil ineffective iu ejecting massive IMDIIs.," Naturally, this implies that if $> 30 M_\odot$ black holes are extremely rare in primordial globular clusters, the retention fraction increases dramatically \citep{2001ApJ...554..548F}, making gravitational recoil ineffective in ejecting massive IMBHs."
+ Figure 6 shows that it is easier to retain an IMDITof less than LOOOAL. with the shallower Belezvuski black hole mass function. with only 30% ejected.," Figure \ref{fig:retainbhmf} shows that it is easier to retain an IMBHof less than $1000 M_\odot$ with the shallower Belczynski black hole mass function, with only $30\%$ ejected."
+ Recall that cach formation chamnel is expected to take place in a particular proto-elobular cluster civiromment single runaways ought to form in very dense systelus.), Recall that each formation channel is expected to take place in a particular proto-globular cluster environment single runaways ought to form in very dense systems.)
+ Since predictions of pre-core collapse clusters indicate that the densities are ~10 times higher than today. we can use current observations of the central Iunünuositv deusitv (?)— to infer the iuifial concditious of the elobular cluster.," Since predictions of pre-core collapse clusters indicate that the densities are $\sim 10$ times higher than today, we can use current observations of the central luminosity density \citep{Harris:96globs} to infer the initial conditions of the globular cluster."
+ Then. we can estimate on a case- basis what IMDII formation channel melt have been appropriate for cach globular cluster.," Then, we can estimate on a case-by-case basis what IMBH formation channel might have been appropriate for each globular cluster."
+ If we assume that every globular cluster forms au IEMDII with a mass that is consisteut with its presumed formation chanucl (see previous section). we can estimate the number of surviving IMDIIs within the Milky. Wav elobular cluster," If we assume that globular cluster forms an IMBH with a mass that is consistent with its presumed formation channel (see previous section), we can estimate the number of surviving IMBHs within the Milky Way globular cluster"
+and in column 4 the apparent visual magnitude.,and in column 4 the apparent visual magnitude.
+ Effective temperatures (Ty) and surface gravities (logg) deduced from the αυ photometry are collected in columns 5 and 6., Effective temperatures $_{\rm{eff}}$ ) and surface gravities $\log g$ ) deduced from the $uvby\beta$ photometry are collected in columns 5 and 6.
+ The derived projected rotational velocities and. microturbulence velocities are collected in columns 7 and 8., The derived projected rotational velocities and microturbulence velocities are collected in columns 7 and 8.
+ Comments about the membership. binarity and pulsation appear in the last column.," Comments about the membership, binarity and pulsation appear in the last column."
+" The rotatioral velocities of these stars were found to range from 7.5 kmss! to 200 kmss7!. six of the A stars are rotating with a v,sini larger than 100 kmss7!."," The rotational velocities of these stars were found to range from 7.5 $^{-1}$ to 200 $^{-1}$, six of the A stars are rotating with a $v_{e}\sin i$ larger than 100 $^{-1}$."
+ According to the CCDM catalogue (Dommanget Nys 1995). HD23387 is the primary in a spectroscopic biary.," According to the CCDM catalogue (Dommanget Nys 1995), HD23387 is the primary in a spectroscopic binary."
+ [ts companion located at 0.3 aresec. has a visual magnitude V29 and should contribute about of the light in the spectrum.," Its companion located at 0.3 arcsec, has a visual magnitude V=9 and should contribute about of the light in the spectrum."
+ The spectral type of this companion is unknown., The spectral type of this companion is unknown.
+ Careful inspection of the spectrum of HD23387 does not reveal lines which could be attributed to à companion of a later spectral type than that of the primary., Careful inspection of the spectrum of HD23387 does not reveal lines which could be attributed to a companion of a later spectral type than that of the primary.
+ The abundances deduced for HD23387 should be taken with The reduction of ELODIE's spectra was fully explained in Paper I and follows the method of Erspamer&North(2002).., The abundances deduced for HD23387 should be taken with The reduction of ELODIE's spectra was fully explained in Paper I and follows the method of \cite{2002A&A...383..227E}.
+ A similar reduction procedure was applied to SOPHIE spectra., A similar reduction procedure was applied to SOPHIE spectra.
+ The correction for scattered light was found to be about in the blue part of the spectrum and less elsewhere., The correction for scattered light was found to be about in the blue part of the spectrum and less elsewhere.
+ The most appropriate method to derive abundances Is spectrum synthesis as several of the investigated stars are fast rotators., The most appropriate method to derive abundances is spectrum synthesis as several of the investigated stars are fast rotators.
+It is thus clear that if radial miotious of some kind are present in the torus. they can iuduce vertical oscillations. and both of these occur at the epicvlic frequencics for test particles in nearly circular orbits.,"It is thus clear that if radial motions of some kind are present in the torus, they can induce vertical oscillations, and both of these occur at the epicylic frequencies for test particles in nearly circular orbits."
+ We now consider a situation in which there is a periodic perturbation at work. which is external to the disk.," We now consider a situation in which there is a periodic perturbation at work, which is external to the disk."
+ The pulsar at the ceuter of the eravitational potential well is clearly an example of this. and one would expect it to have an effect on the disk.," The pulsar at the center of the gravitational potential well is clearly an example of this, and one would expect it to have an effect on the disk."
+" The magnetic field of the pulsar. or sole deformation on its surface. can perturb the disk at intervals given by the inverse of the spin xiod. AT=l/r, Gu the case of SAN Jlsds.-3658. Ότι. Uz)."," The magnetic field of the pulsar, or some deformation on its surface, can perturb the disk at intervals given by the inverse of the spin period, $\Delta
+T=1/\nu_{s}$ (in the case of SAX J1808.4-3658, $\nu_{s}$ =401 Hz)."
+ We have cousidered then a corcine in the radial direction. which manifests itself through a small. radia acceleration. the naegnitude of which is shown in Fie. 3..," We have considered then a forcing in the radial direction, which manifests itself through a small, radial acceleration, the magnitude of which is shown in Fig. \ref{forcingamp}."
+ It is not o»urelv smusoidal but repeats at a fixed interval AT. as given above.," It is not purely sinusoidal but repeats at a fixed interval $\Delta T$, as given above."
+ The adopted profile of this oerturbation deserves some comment: we have used it to mimic the passage of a brief (compared with the repetition time). but fairly stronger than average. disturbance in the accretiondisk (c.g.. the corresponding polar magnetic field of the pulsar. sWweeopiug aroule the disk).," The adopted profile of this perturbation deserves some comment: we have used it to mimic the passage of a brief (compared with the repetition time), but fairly stronger than average, disturbance in the accretion disk (e.g., the corresponding polar magnetic field of the pulsar, sweeping around the disk)."
+ Diffoyeu shapes for this pulse will be explored elseschere., Different shapes for this pulse will be explored elsewhere.
+ The radial forciug obviously induces a radial oscillation of the torus., The radial forcing obviously induces a radial oscillation of the torus.
+ Because of pressure coupling. a vertical oscillation 1s again apparent. albeit at a reduced maeuitude (compared with the radial amplitude).," Because of pressure coupling, a vertical oscillation is again apparent, albeit at a reduced magnitude (compared with the radial amplitude)."
+ The motion of the ceuter of the torus (as defined above) can be Fourieranalyzed to extract the relevant frequencies., The motion of the center of the torus (as defined above) can be Fourier--analyzed to extract the relevant frequencies.
+ We fud that the radial aud vertical motions occur priuarily at the local radial aud epicvlic frequencies. Ky aud Cy respectively. at ry.," We find that the radial and vertical motions occur primarily at the local radial and epicylic frequencies, $\kappa_{0}$ and $\zeta_{0}$ respectively, at $r_{0}$ ."
+" For a torus orbiting a mass AM= 1.28M.. aud ceuter at ry=12.25 δισ, sg=500 IIz aud Cy=TUO TIz."," For a torus orbiting a mass $M=1.38$ $_{\odot}$, and center at $r_{0}=12.25M=24.8\,$ km, $\kappa_{0}=500$ Hz and $\zeta_{0}=700$ Hz."
+ The question now is: what effect. if amy. does the repetition time AT of the perturbation lave ou tle power of oscillatory motion iuduced in the torus?," The question now is: what effect, if any, does the repetition time $\Delta T$ of the perturbation have on the power of oscillatory motion induced in the torus?"
+ The radial motion is always driven at a relatively hieh amplitude. simply because the perturbation is itself applied as a radial acceleration.," The radial motion is always driven at a relatively high amplitude, simply because the perturbation is itself applied as a radial acceleration."
+ We fiud. however. that the power of the vertical motions varies ereatlv as AT is altered.," We find, however, that the power of the vertical motions varies greatly as $\Delta T$ is altered."
+ We have kept the initial position of the torus fixed at ry=12.2547 and performed over a dozen sinulations. differing ouly in the value of νε. and covering a range 100<(Iz)«600.," We have kept the initial position of the torus fixed at $r_{0}=12.25M$ and performed over a dozen simulations, differing only in the value of $\nu_{s}$, and covering a range $100<\nu_{s} \mbox{(Hz)}<600$."
+ For each one of these. we have computed the peak power Dag. du the vertical oscillation at cy.," For each one of these, we have computed the peak power $P_{M,z}$ in the vertical oscillation at $\zeta_{0}$."
+ Since one could equivalently keep 74; fixed (as is actually the case in nature) aud vary ry Gvhich would alter the difference GyKy). iu Fig.," Since one could equivalently keep $\nu_{s}$ fixed (as is actually the case in nature) and vary $r_{0}$ (which would alter the difference $\zeta_{0}-\kappa_{0}$ ), in Fig."
+ { we show Pay. asa function of the ratio (ο&g).," \ref{vertpower} we show $P_{M,z}$ as a function of the ratio $\nu_{s}/(\zeta_{0}-\kappa_{0})$."
+ The response of the torus is clearly greatest when νε=2(cy5g). as was observed in the single instance for SAN Jisos.3658 when two OPO peaks iu the kIIz Paneer were seen.," The response of the torus is clearly greatest when $\nu_{s}=2 (\zeta_{0}-\kappa_{0})$, as was observed in the single instance for SAX J1808.4-3658 when two QPO peaks in the kHz range were seen."
+" We note that there is also a strong response whenv, aud ορA#y are ia Lil or 3:2 correspondence.", We note that there is also a strong response when $\nu_{s}$ and $\zeta_{0}-\kappa_{0}$ are in a 1:1 or 3:2 correspondence.
+ The first of these would allow for the possibility of twin peaks with a separation of 101 Tz in SAN JLsos.3658. while the second would imply a separation of 802/3=267 Uz.," The first of these would allow for the possibility of twin peaks with a separation of 401 Hz in SAX J1808.4-3658, while the second would imply a separation of $802/3=267$ Hz."
+ The first option would occur at οΞ105 1 Tz. 592653 Tz (see also Whiguialketal. (2003))). auc ry =9.75\I=19.7 kan for 1.98M....," The first option would occur at $\zeta_{0}$ =1054 Hz, $\kappa_{0}$ =653 Hz (see also \citet{kakls03}) ), and $r_{0}$ =9.75M=19.7 km for $_{\odot}$."
+ The second would imply cy=s832 Tz and sg565 Iz at a radius cj—11.13M2—22.5 lan for the same mass., The second would imply $\zeta_{0}$ =832 Hz and $\kappa_{0}$ =565 Hz at a radius $r_{0}$ =11.13M=22.5 km for the same mass.
+ Neither has heen observed as vet., Neither has been observed as yet.
+ The appearance of twin kIIZ QPOs iu the millisecond pulsar SAN JLS0s.1-3658. with a separation consistent with half the known spin frequency. of the pulsar. stronely indicates that a nonlinear resonance is at work. coupling the spin to vibrational modes in the disk.," The appearance of twin kHZ QPOs in the millisecond pulsar SAX J1808.4-3658, with a separation consistent with half the known spin frequency of the pulsar, strongly indicates that a nonlinear resonance is at work, coupling the spin to vibrational modes in the disk."
+ Using a simple lvdvodvuamical model. we ideutify these modes with the radial aud epicyclic oscillations of fluid elements slightly displaced from exact circular orbits.," Using a simple hydrodynamical model, we identify these modes with the radial and epicyclic oscillations of fluid elements slightly displaced from exact circular orbits."
+ In this respect. the model is crucially dependent on the effects of stroug gravity. to break the degeneracy between the orbital and epicyclic frequencies preseut in the Newtouiau regime.," In this respect, the model is crucially dependent on the effects of strong gravity, to break the degeneracy between the orbital and epicyclic frequencies present in the Newtonian regime."
+ Under the unique assumption that the pulsar provides a periodic drivingradial force to a sleuder torus in orbit GQvlich we cousider as a standin for a density cuhancement in the accretion disk: see 32). we show that the response of the torus," Under the unique assumption that the pulsar provides a periodic drivingradial force to a slender torus in orbit (which we consider as a stand–in for a density enhancement in the accretion disk; see \ref{response}) ), we show that the response of the torus"
+Space Office Grant No.,Space Office Grant No.
+ URK09350 and the “Lendiillet’ program of the Hungarian Academy of Sciences., URK09350 and the `Lendüllet' program of the Hungarian Academy of Sciences.
+ KK acknowledges the support of Austrian FWF projects T359 and PI19962., KK acknowledges the support of Austrian FWF projects T359 and P19962.
+ The authors gratefully acknowledge the entireKepler team. whose outstanding efforts have made these results possible.," The authors gratefully acknowledge the entire team, whose outstanding efforts have made these results possible."
+and the period is There are two sources for Fluctuations in the potential. those caused. by the protons. and. those. caused by the electrons.,"and the period is There are two sources for fluctuations in the potential, those caused by the protons and those caused by the electrons."
+ The typical time scales of the Ductuations are where υπ.7 is the mean interparticle distance and Aion is the number density of ions., The typical time scales of the fluctuations are where $\langle r_{s} \rangle=1/(4 \pi n_{{\rm ion}}/3)^{1/3}$ is the mean interparticle distance and $n_{{\rm ion}}$ is the number density of ions.
+" ey, is the relevant thermal velocity.", $v_{th}$ is the relevant thermal velocity.
+ Phe number of particles in the Debve sphere is given by: is the Debve radius for this mixture., The number of particles in the Debye sphere is given by: where is the Debye radius for this mixture.
+) is the number density of specie j with charge Z;., $n_{j}$ is the number density of specie $j$ with charge $Z_{j}$.
+ The above expression for the Debye radius assumes that both the electrons and ions contribute to the supposedly: DII potential., The above expression for the Debye radius assumes that both the electrons and ions contribute to the supposedly DH potential.
+ For simplicity we assume ο.-..2:3 and obtain. that for⋅ n—10726ácm’. TH15.I0' Ix anda pure Hydrogen plasma. Note that when some of the ions are Lelium ions (in he core of the present Sun about half the ions are He). Nip decreases even more.," For simplicity we assume $N_{D}^{1/2}
+\approx 3$ and obtain that for $n=10^{26}\# /{\rm cm}^3$, ${\rm
+T}=1.5\times 10^{7}$ K and a pure Hydrogen plasma, Note that when some of the ions are Helium ions (in the core of the present Sun about half the ions are He), $N_{D}$ decreases even more."
+ What is the implication of this result?, What is the implication of this result?
+ The Huctuations due to the electrons are of about the same imeseale as the period of the electrons in the Ix-shell in a single ion in vacuum. while the protons in the plasma have a much longer time scale.," The fluctuations due to the electrons are of about the same timescale as the period of the electrons in the K-shell in a single ion in vacuum, while the protons in the plasma have a much longer time scale."
+ Hence. it is not justified to treat he contribution of the electrons to the DII potential (felt w the electron) as a smooth potential in time.," Hence, it is not justified to treat the contribution of the electrons to the DH potential (felt by the electron) as a smooth potential in time."
+ On the other ancl. since there are only few protons in the Debve radius heir contribution to the potential is smooth in time but not in space and certainly not spherical.," On the other hand, since there are only few protons in the Debye radius their contribution to the potential is smooth in time but not in space and certainly not spherical."
+ Additional arguments hat question the validitv of the potential are given by 2.., Additional arguments that question the validity of the potential are given by \citet{john92}.
+ In what follows. we assume that all potentials are temporally smooth and spatially spherical.," In what follows, we assume that all potentials are temporally smooth and spatially spherical."
+ Next. we discuss the ionization state of Bet in the core of the Sun - the classical wav.," Next, we discuss the ionization state of ${\rm Be}^{7}$ in the core of the Sun - the classical way."
+" I one adopts the Saha equation. ignoring screening and the excited energy. levels (thus including only the ground. states in the partition functions). then the probabilities fj ancl f» that one or two Ix-shell electrons are associated with any given De! nucleus are given by (Hx867) where Hore x,=216.6eV is the forth ionization potential of he Be! atom. v»=153.1eV is the third ionization potential of the ο atom. and AY=x4.Χο63.5eV."," If one adopts the Saha equation, ignoring screening and the excited energy levels (thus including only the ground states in the partition functions), then the probabilities $f_{1}$ and $f_{2}$ that one or two K-shell electrons are associated with any given ${\rm Be}^{7}$ nucleus are given by (IKS67) where Here $\chi_{1}=216.6 {\rm ~eV}$ is the forth ionization potential of the ${\rm Be}^{7}$ atom, $\chi_{2}=153.1 {\rm ~eV}$ is the third ionization potential of the ${\rm Be}^{7}$ atom, and $\Delta \chi =
+ \chi_{1} - \chi_{2} = 63.5 {\rm ~eV}$."
+ Phese values correspond. to the limit of vanishing plasma density., These values correspond to the limit of vanishing plasma density.
+" n, is he number density of the free electrons. most of which are contributed by IEvdrogen and Helium and. are independent of the state of trace elements like Beryllium."," $n_{e}$ is the number density of the free electrons, most of which are contributed by Hydrogen and Helium and are independent of the state of trace elements like Beryllium."
+ Thus. Εν can » treated as fixed.," Thus, $n_{e}$ can be treated as fixed."
+" The application of the Saha equation in the above Form o the core of the Sun (p=158gem7. T=L57«10kx. Xom0.36 Z=0.02) vields f,=0.320 and |f.0.038. implving that the Bet keeps its last electron for about a hird of the time."," The application of the Saha equation in the above form to the core of the Sun $\rho =158 {\rm ~g~cm^{-3}}$, ${\rm T} =1.57 \times
+10^{7}{\rm K}$, $X=0.36$ $Z=0.02$ ) yields $f_{1}=0.320$ and $f_{2}=0.038$, implying that the ${\rm Be}^{7}$ keeps its last electron for about a third of the time."
+ As the relevant ions are in a plasma. the traditional procedure to correct. for the plasma cllect is to replace he pure Coulomb potential with a DII one (2)..," As the relevant ions are in a plasma, the traditional procedure to correct for the plasma effect is to replace the pure Coulomb potential with a DH one \citep{RGH70}."
+ This is or example the procedure Uss67 evaluated. the plasma corrections for the energy levels of Be! in the solar core.," This is for example the procedure IKS67 evaluated the plasma corrections for the energy levels of ${\rm
+Be}^{7}$ in the solar core."
+ The Saha equation in the above form ignores clectron degeneracy. exchange cliccts and pressure ionization.," The Saha equation in the above form ignores electron degeneracy, exchange effects and pressure ionization."
+ The electron degeneracy introduces a small correction under the conditions prevailing in the Sun., The electron degeneracy introduces a small correction under the conditions prevailing in the Sun.
+ As we will shortly demonstrate. exchange ancl pressure ionization are significantly more important.," As we will shortly demonstrate, exchange and pressure ionization are significantly more important."
+ In what follows we do assume in spite of the previous reservations. à smooth static DIL potential contributed by the electrons and ions.," In what follows we do assume in spite of the previous reservations, a smooth static DH potential contributed by the electrons and ions."
+ Moreover. we assume it to be relevant in a statistical sense only.," Moreover, we assume it to be relevant in a statistical sense only."
+ After performing the above estimate. LIXS67 turned to evaluate the ground. state of the Z=4 ion assuming a smooth DII screened. potential in which both the protons and the electrons are taken into account.," After performing the above estimate, IKS67 turned to evaluate the ground state of the $Z=4$ ion assuming a smooth DH screened potential in which both the protons and the electrons are taken into account."
+ We ignored. the questions raised in the previous section concerning the validity of the potential for our particular purpose here (ionization in the core of the Sun). repeated their calculation and confirmed. their results with respect to the Lvdrogen like ion with Z—4.," We ignored the questions raised in the previous section concerning the validity of the potential for our particular purpose here (ionization in the core of the Sun), repeated their calculation and confirmed their results with respect to the Hydrogen like ion with $Z=4$."
+ ?. calculated the bound states of static screened. coulomb potential and formulated their results in terms of the screening length., \citet{RGH70} calculated the bound states of static screened coulomb potential and formulated their results in terms of the screening length.
+ We also repeated their results for the Ts state and the results are shown in fig. 1.., We also repeated their results for the $1s$ state and the results are shown in fig. \ref{fig:screenE}.
+ Clearly. as the Debye length. approaches the Bohr radius of ions with charge Z. there are no more bound states.," Clearly, as the Debye length approaches the Bohr radius of ions with charge Z, there are no more bound states."
+ The boundary condition on the wave function in this case is c(rx)=0., The boundary condition on the wave function in this case is $\psi(r \rightarrow \infty ) =0$.
+ The calculation. of the plasma effects on the triply ionized De* ion is more complicated because of the partial screening of the nucleus by the bound electron., The calculation of the plasma effects on the triply ionized ${\rm Be}^{7}$ ion is more complicated because of the partial screening of the nucleus by the bound electron.
+ To overcome this problem. we used the following approximate method.," To overcome this problem, we used the following approximate method."
+ We Iooked for the eigenvalue in the low density limit and searched. for the effective. charge that will reproduce the, We looked for the eigenvalue in the low density limit and searched for the effective charge that will reproduce the
+ The orderofthe QCDtrausition witha simallcomparedlatent,"after the transition completed, is This nucleation distance sets the spatial scale for baryon number inhomogeneities."
+ heat.dimensional compared othe bagl.. model.We aud athatsmall surfa, Lattice simulations imply that in real-world QCD the energy density must change very rapidly in a narrow temperature interval.
+ceOCD rausitionension. isoffirstto orderaud that arguiueuts hevalues from ," This can be seen from the microscopic sound speed in the quark phase, $c_s \equiv (\partial p/\partial \varepsilon)_{\rm S}^{1/2}$."
+assuuequenched latticethe OCD appropriately, Lattice QCD indicates that $3 c_{\rm s}^2(T_{\rm c}) = {\cal O}(0.1)$ \cite{Latticecs2}.
+ by the. ofdegrees offreedom) typice Ulfor (scaled.physical QCD transition. Baseduuuber these values hoare bubble he wicleation small A.," Thus, the cosmological time-temperature relation is strongly modified already before the nucleations,due to where $t_H \equiv 1/H = (3 M_{\rm pl}^2/8\pi\varepsilon_{\rm q})^{1/2}$ with $\varepsilon_{\rm q}$ being the energy density in the quark phase."
+ oji1 Tif Ti.aux101 HOLCUCOusbubble meleation distance.a supercooling.diye Lem.woul =|follow?.. Theactial.nucleatio, This behavior of the sound speed increases the nucleation distance because of the proportionality $\Delta t_{\rm nuc}\propto 1/[3 c_{\rm s}^2 (T_f)]$ \cite{Ignatius}.
+nand ati the earlyargueUniversetheassumptioninevitable «ensitvlolnuogoneous perturbations from iufiation, In the thin-wall approximation the nucleation action has the following explicit expression: for small supercooling.
+ or from other seeds forbv structure Oraion. Those fuctuations ducdi‘sity and have beenmicased byCODE!Q wavean éT/T ~10 temperatureDThe effectoftheQC 'D transitionon," Assuming further that $c_{\rm s}$ does not change very much during supercooling, the following relation holds for the supercooling and nucleation scales: Here we denote by $\Delta$ a relative (dimensionless) temperature interval and by $\Delta t$ a dimensionful time interval."
+ ¢lesity perturbatio aiupliti lisD.G and," $\bar{S} \equiv S(T_f)$ is the critical nucleation action, $\bar{S} = {\cal O}(100)$."
+ eravitational inpuntiesare present.the τοςαμάνis homogeneous uucleation. Theprobability to nucleatea," Surface tension and latent heat are provided by lattice simulations with quenched QCD only, giving the values $\sigma = 0.015
+T_c^3$, $l = 1.4 T_c^4$ \cite{Iwasaki}."
+ bubleofthe new phaseper timeaud vo Tueis approximated by T= Tlexp| , Scaling the latent heat for the physical QCD leads us to take $l = 3 T_c^4$ .
+"ΤΗ.The cooling of theU uverse. The durationofthe mucleatio1period. Δήμο,is found to be *? Ati EL (1) dsdff,"," With these values for the latent heat and surface tension, the amount of supercooling is $\Delta_{\rm sc} = 2.3 \times 10^{-4}$."
+ time tyis definedas , From Eq. \ref{Seq}) )
+the uonment when the fractionof space where n, it follows that $\Delta_{\rm nuc} = 1.5 \times 10^{-6}$ .
+ucleatious stillcontinue equals1/6.The," Substituting $3 c_{\rm s}^2 = 0.1$ into Eq. \ref{Tteq}) ),"
+heatflow preceding the deflaeration fronts rcheats ofthe Universe.We denote, we find $\Delta t_{\rm nuc} = 1.5 \times 10^{-5} t_{\rm H}$ for the duration ofthe nucleation period.
+ bw Cua hee ffectivespeedby which released latentheat propaga esm sufficient, The nucleation distance depends on the unknown velocity $v_{\rm heat}$ in Eq. \ref{dnuceq}) ).
+ amoutst oslut down uucleations. Iu practice.CacD Όμως € ey. Where ομως theve ocity of the deflagrationfrout andοςds," With the value 0.1 for $v_{\rm heat}$, the nucleation distance $d_{\rm nuc,hom}$ would have the value $2.9 \times 10^{-6} d_{\rm H}$."
+ the soundspeed 9. Iu the unlikelycaseof detonations(year Slot Idbe replacedby the v," One should take these values with caution, due tolarge uncertaintiesin $l$ and $\sigma$ ."
+elocitv oft1ο phaseboundary1 n all expressions thatfo low. Theien « listance between uucleationce iters.measured inuuediatelv," As our reference set of parameters, we take: $\Delta_{\rm sc} = 10^{-4}$ , $\Delta_{\rm nuc} = 
+10^{-6}$ , $\Delta t_{\rm nuc} = 10^{-5} t_H$ ."
+cixk.,disk.
+ A detection of the Zeemau effect in a disk has previously been reported by Donati ((2005) in the case of FU Ori., A detection of the Zeeman effect in a disk has previously been reported by Donati (2005) in the case of FU Ori.
+ This is thought to be an accretion disk with a disk wind. as evidenced by some lines showing P. Cyveui absorption.," This is thought to be an accretion disk with a disk wind, as evidenced by some lines showing P Cygni absorption."
+ However. he cimeular polarization profile is argued as cine associated with the iunermost region of he rotating disk where the field is of kilogauss streneths.," However, the circular polarization profile is argued as being associated with the innermost region of the rotating disk where the field is of kilogauss strengths."
+" The detection iuplics a net magnetic Hux per spectral resolution clement. and thus sets limuts on the turnover (or ""tauglednuess) cheth scale for the disk field. if iudeed the MBI is operating in this case."," The detection implies a net magnetic flux per spectral resolution element, and thus sets limits on the turnover (or “tangledness”) length scale for the disk field, if indeed the MRI is operating in this case."
+ It is interesting to consider signals that could result with the Tale effect., It is interesting to consider signals that could result with the Hanle effect.
+ The Iaule effect can operate in regions where magnetic ficlds are “tangled” or randomized., The Hanle effect can operate in regions where magnetic fields are “tangled” or randomized.
+ This meaus that spatial averages of (HB tend toward zero although 325 does not. such as is the case for the MRI mechanisin.," This means that spatial averages of $\langle \vec{B} \rangle$ tend toward zero although $\langle B^2 \rangle$ does not, such as is the case for the MRI mechanism."
+ Au exteusive literature exists for the ILuie effect with random fields iu applications to solar studies (e.g... Frisch 22009. aud references therein).," An extensive literature exists for the Hanle effect with random fields in applications to solar studies (e.g., Frisch 2009, and references therein)."
+" Tere I simply ""aut to outline some of the limiting behavior iu applications to clisks.", Here I simply want to outline some of the limiting behavior in applications to disks.
+ Consider a I&epleriau disk that contains evervcliere a truly raudonmized magnetic field., Consider a Keplerian disk that contains everywhere a truly randomized magnetic field.
+ If the field. is weak at all locations. mcaning that D«Di. then of. course a line∙ profile⋅ results as ↗∙iu tle case of. no ITaule effect.," If the field is weak at all locations, meaning that $B \ll \BHan$, then of course a line profile results as in the case of no Hanle effect."
+. ) itME the fiMfield is1 strong.. strong.suchsuc that: thee deuserdenser regionsBut of the lddisk are largely saturated. then the behavior is much different.," But if the field is strong, such that the denser regions of the disk are largely saturated, then the behavior is much different."
+ With different field orieutatious. oue expects that wwill ∙∙vield a null profile.∙ by svuuuetry considerations.," With different field orientations, one expects that will yield a null profile, by symmetry considerations."
+. I profile NONE.is thedifferent.," However, the profile is different."
+US As a specific example. consider the resulting line from an edge-on disk.," As a specific example, consider the resulting line from an edge-on disk."
+ Without a feld. the »olarization at line center would be zero. owiug to orward scattering of unpolarized starlight.," Without a field, the polarization at line center would be zero, owing to forward scattering of unpolarized starlight."
+ With a randomized feld. the polarization will still teud OWAward ZOYOzero ati lineHο cougux.," With a randomized field, the polarization will still tend toward zero at line center."
+ In the poiut star laut. sole net poapolarization is expected] to survive in the ine wines.," In the point star limit, some net polarization is expected to survive in the line wings."
+ This polarization will be significautlv reduced in comparison to the zero field case., This polarization will be significantly reduced in comparison to the zero field case.
+ If Ly were unitw. the polarization at the line wings would be.. since the scattering geonmetrv is 907.," If $E_1$ were unity, the polarization at the line wings would be, since the scattering geometry is $90^\circ$."
+ Tn analogy with cousideratious of scattering polarization off the solar limb. a reduction iu polarization by a factor of 5 for isotropically distributed felds should be expected (see Stento 1982).," In analogy with considerations of scattering polarization off the solar limb, a reduction in polarization by a factor of 5 for isotropically distributed fields should be expected (see Stenflo 1982)."
+ Now consider the introduction of a sustained oroidal field Compoucut., Now consider the introduction of a sustained toroidal field component.
+ Of course toroidal fields were considered in the previous section., Of course toroidal fields were considered in the previous section.
+ Now rowever. the toroidal field has polarity flips within he disk B.. is alternately ew or cow at essentially random points within the isovelocity zones.," Now however, the toroidal field has polarity flips within the disk – $B_\varphi$ is alternately cw or ccw at essentially random points within the isovelocity zones."
+ What lis means is that there is a sien chauge in the direction of Larmor precession iu the classical victure of a harmonic oscillator., What this means is that there is a sign change in the direction of Larmor precession in the classical picture of a harmonic oscillator.
+ The effect of this is to drive the ssignal to zero faster than if the toroidal field had one seuse of polarity., The effect of this is to drive the signal to zero faster than if the toroidal field had one sense of polarity.
+ In the saturated limit. the pprofile d$.unchanged. lecause the surviving polarized signal does uot depend ou polarity at all.," In the saturated limit, the profile is, because the surviving polarized signal does not depend on polarity at all."
+ This is quite different from the Zecnian effect that is sensitive to the net maguetic fux., This is quite different from the Zeeman effect that is sensitive to the net magnetic flux.
+" For the Zocinan effect. the circular polarization will be suppressed when D, switches polarity on sumiall spatial scales."," For the Zeeman effect, the circular polarization will be suppressed when $B_\varphi$ switches polarity on small spatial scales."
+ Polarized5⋅ line⋅ profile⋅ shapes from. magnetized. Keplerian-. disks. have been calculated under a uunber of simplifving assumptions: the disk is ecometrically thiu: the scatteriug lines are optically thin: primarily simple fields were cousidered (axialP or. toroidal):⋅∙. aud no account was⇁↴ taken↽ of photospheric absorption lines., Polarized line profile shapes from magnetized Keplerian disks have been calculated under a number of simplifying assumptions: the disk is geometrically thin; the scattering lines are optically thin; primarily simple fields were considered (axial or toroidal); and no account was taken of photospheric absorption lines.
+ On the other haud. the 110del profiles⋅ do inchide⋅ fuite⋅⋅ source devolarization aud the effects of stellar occultation.," On the other hand, the model profiles do include finite source depolarization and the effects of stellar occultation."
+" The presentationof results focused on the polarimetric ""efficiencies"" aand wwith a description of how to identify the occurrence of the ILaule effect in scattering lines from disks.", The presentation of results focused on the polarimetric “efficiencies” and with a description of how to identify the occurrence of the Hanle effect in scattering lines from disks.
+ Tn addition. a discussion was presented for the Taule effect from a magnetized disk iu which the ΑΠΗΤ mechanisma is operating.," In addition, a discussion was presented for the Hanle effect from a magnetized disk in which the MRI mechanism is operating."
+ One of the main couchlisious from this work iz hat axial aud toroidal fields in disks mre casily distinguishable through an analysis of mfigures forn the run ofB Stokes1 polarizations.H across ine: profiles., One of the main conclusions from this work is that axial and toroidal fields in disks are easily distinguishable through an analysis of figures for the run of Stokes polarizations across line profiles.
+. Although↜ a yprofile does exist even without the TWaule effect. owing to stellar occultation. its amplitude is quite siuall.," Although a profile does exist even without the Hanle effect, owing to stellar occultation, its amplitude is quite small."
+ The strongest »profiles result when auch of the iuner disk. where uost of the scattered light is produced. has values |: BíDqya of order a few.," The strongest profiles result when much of the inner disk, where most of the scattered light is produced, has values of $B/\BHan$ of order a few."
+ If the iuuer disk is uostlv in the saturated Limit of the ITaule effect.," If the inner disk is mostly in the saturated limit of the Hanle effect,"
+thanTEMPO?2.,than.
+. The accuracy of the simplified timing solution was 300 ns., The accuracy of the simplified timing solution was $300$ ns.
+ Event times were transfered to the solar system barycenter. having corrected for radio dispersion and pulsar proper motion.," Event times were transfered to the solar system barycenter, having corrected for radio dispersion and pulsar proper motion."
+ Times of arrival were phase-folded based on the efpp/tase--compatible version of the ephemeris. yielding the phase histogram in Figure 6..," Times of arrival were phase-folded based on the -compatible version of the ephemeris, yielding the phase histogram in Figure \ref{Parkes_0437}."
+ The mean phase calculated with the LAT software is delayed from the mean value by 0.32 us. resulting in a validation of the code below the jes level.," The mean phase calculated with the LAT software is delayed from the mean value by $0.32$ $\mu$ s, resulting in a validation of the code below the $\mu$ s level."
+ We have motivated and described the large timing campaign that is underway for the, We have motivated and described the large timing campaign that is underway for the
+distortion of the nuages of background galaxies] is eenerallv consistent with the virial discrepancy to within a factor of two or three (Wu Fang 1997. Allen 1998).,"distortion of the images of background galaxies) is generally consistent with the virial discrepancy to within a factor of two or three (Wu Fang 1997, Allen 1998)."
+" Therefore auv relativistic extention of MOND which preserves the relation between the weak field force aud the deflection of light νο, 0=2ομ ⋅ ↙⋅−∕∙↙∕⊥↙∕∣⋟∖∏∐⋜↧↕↴∖↴∪⋜↧↸⊳↸⊳∪∏∐↑ for the lensing discrepaucy. (an example of such a theory ↕↴∖↴∶↴∙⊾↕↖↽↸∖∐↕∐≋⋜∐∐∐∖↥∷∖∷↕"," Therefore any relativistic extention of MOND which preserves the relation between the weak field force and the deflection of light (i.e., $\theta = 2/c^2\int{g_\perp dl}$ ) will also account for the lensing discrepancy (an example of such a theory is given in Sanders, 1997)."
+∩∩⊤⋝∙ Some clusters also act as strong lenses: L6.. multiple nuages of background sources are formed by the ceutral regions of the clusters.," Some clusters also act as strong lenses; i.e., multiple images of background sources are formed by the central regions of the clusters."
+ The critical surface density required for strong lensing is where Fis a diueusionless fiction of the leus aud source redshifts which depeuds upon the cosmological model(Blandford Nvaravan 1992): typically for clusters and background sources at cosiiological distances Fz10., The critical surface density required for strong lensing is where F is a dimensionless function of the lens and source redshifts which depends upon the cosmological model (Blandford Narayan 1992); typically for clusters and background sources at cosmological distances $F\approx 10$.
+" Modified dynamics applies in the Πιτ of low accelerations or. equivaleutly. at surface deusities below a value of Xa,zxαρ CMilgeroimi 198232)."," Modified dynamics applies in the limit of low accelerations or, equivalently, at surface densities below a value of $\Sigma_M \approx a_o/G$ (Milgrom 1983a)."
+" Since. MOL it is found that a,zz«LL,ο, this iuplies that X,zmMap mathat is to sav. the critical surface density for strong is always ereater than the upper limit for MOND phenomenoloey.regine."," Since, observationally it is found that $a_o \approx cH_o/6$, this implies that $\Sigma_c \approx 6\Sigma_M$; that is to say, the critical surface density for strong lensing is always greater than the upper limit for MOND phenomenology."
+" Strong lensing observed in clusters typically requires a total projected mass iu the ner 100-200 Ipc in excess of 1011Af. which is evideutlv not preseut iu the formu of hot eas,", Strong lensing observed in clusters typically requires a total projected mass in the inner 100-200 kpc in excess of $10^{14} M_\odot$ which is evidently not present in the form of hot gas.
+ ITitherto undetected matter does seen to be necessary in the cores of rich clusters which exhibit strong leusiug. even with modified dvuamiucs (see also \Gleroi 1996].," Hitherto undetected matter does seem to be necessary in the cores of rich clusters which exhibit strong lensing, even with modified dynamics (see also Milgrom 1996)."
+ This may be taken as a failure or as a prediction., This may be taken as a failure or as a prediction.
+ That the tally of ordinary barvouic matter lay not vet be complete is suggested by the observations. in several clusters. of diffuse star light (Fereuson et al.," That the tally of ordinary baryonic matter may not yet be complete is suggested by the observations, in several clusters, of diffuse star light (Ferguson et al."
+ 1995) aud ultraviolet ciission apparently from wari clouds (Mittaz et al., 1998) and ultraviolet emission apparently from warm clouds (Mittaz et al.
+ 1998)., 1998).
+ Moreover. many X-ray clusters show evidence for cooling flows at somedisappears level {Sarazin LOSS).," Moreover, many X-ray clusters show evidence for cooling flows at some level (Sarazin 1988)."
+ The gas cools. flows inward aux iufo sole. as vet. undetectable form.," The gas cools, flows inward and disappears into some, as yet, undetectable form."
+ The mass deposition rates are generally too low to be dynamically significant at the xeseut epoch. but this may not have always been the Case.," The mass deposition rates are generally too low to be dynamically significant at the present epoch, but this may not have always been the case."
+ lu ΠΑΝ. MOND oOeoes a longC» wav in resolvingC» je virial discrepancy in clusters.," In summary, MOND goes a long way in resolving the virial discrepancy in clusters."
+ T cannot resolve the strous lensing discrepancy iu those clusters where this phenomenon is observed. but this leads to the prediction lat anere barvonic matter ids present anc possibly letectable in the cores of rich clusters.," It cannot resolve the strong lensing discrepancy in those clusters where this phenomenon is observed, but this leads to the prediction that more baryonic matter is present and possibly detectable in the cores of rich clusters."
+ Tam grateful to M. Milgroii for useful comments ou us work., I am grateful to M. Milgrom for useful comments on this work.
+As iu the uuucrical computations. we have based our opacity law on the work of ΟΠ (thin ico. density 10° ).,"As in the numerical computations, we have based our opacity law on the work of OH (thin ice, density $10^6$ )."
+ For the purposes of our analytic work. a very crude uecewise power law fit with breaks at 10 and LOO thas been used (seeTable B2)).," For the purposes of our analytic work, a very crude piecewise power law fit with breaks at 10 and 400 has been used (seeTable \ref{op_param}) )."
+" We then assume that in a sufficicutly large range of wavelengths around the reals of cach compoucut. the dust absorption efficiency Q, can be approximated by a power-law. (Q,X(AyAye."," We then assume that in a sufficiently large range of wavelengths around the peak of each component, the dust absorption efficiency $Q_\nu$ can be approximated by a power-law, $Q_\nu \propto
+(\lambda_{\rm p}/\lambda)^\alpha$."
+ The values of the parameters Aj and a usec Dy us are isted in Table D2.., The values of the parameters $\lambda_{\rm p}$ and $\alpha$ used by us are listed in Table \ref{op_param}.
+ Let the radiation absorbed (indicated by the subscript va) by dust be concentrated around a frequency 14. aud he radiation enütted (udicated by the subscript πο} by dust be characterized by a frequency 1.," Let the radiation absorbed (indicated by the subscript ”) by dust be concentrated around a frequency $\nu_a$, and the radiation emitted (indicated by the subscript ”) by dust be characterized by a frequency $\nu_e$."
+" We take 1η. and a, equal to the values for absorption iu the FIR (see Table B2)). aud we consider the heating of the dust by each componcut of the radiation feld separately."," We take $\nu_e$ and $\alpha_e$ equal to the values for absorption in the FIR (see Table \ref{op_param}) ), and we consider the heating of the dust by each component of the radiation field separately."
+" With these assmuptions. the mean intensity at the cloud’s center is = |. where 7, is the ceutre-to-edge optical depth of the cloud at frequency μι"," With these assumptions, the mean intensity at the cloud's center is = , where $\tau_a$ is the centre-to-edge optical depth of the cloud at frequency $\nu_a$."
+" Then. defining 3%,=ADTην... 4DAT;/ liv, eq. (1))"," Then, defining $\beta_d\equiv kT_d/h\nu_e$, $\beta_i\equiv
+kT_i/h\nu_a$ , eq. \ref{therm_eq}) )"
+" becomes cQ).INE- where Q, aud Q. are the values of Q, at frequencies pu and 14. respectively. aud dit didt. with g=a,|p3."," becomes )^p _i W_i where $Q_a$ and $Q_e$ are the values of $Q_\nu$ at frequencies $\nu_a$ and $\nu_e$, respectively, and dt _i= dt, with $q=\alpha_a+p+3$."
+ Iu the above. P aud ¢ are the eanuna and Rienuun zeta fuuctious (see Abramovitz Steeuu 1965).," In the above, $\Gamma$ and $\zeta$ are the gamma and Riemann zeta functions (see Abramovitz Stegun 1965)."
+" The inteeral in the equation for B; cau be casily evaluated in the two limiting situations 7,<<1 and TQ1.", The integral in the equation for ${\cal B}_i$ can be easily evaluated in the two limiting situations $\tau_a\ll 1$ and $\tau_a\gg 1$.
+" For 7,<1 the exponoeutial iu B; cau be expanded asa Taylor series. giving. to first order in 7,. Thus. the dust temperature at the ceuter is given. to lowest order. bv andthecloudEs. where 1 C(|I"," For $\tau_a\ll 1$ the exponential in ${\cal B}_i$ can be expanded as a Taylor series, giving, to first order in $\tau_a$, Thus, the dust temperature at the center is given, to lowest order, by C _i }, where )^4 }."
+"rποPleta,li1D0,QQ,UmMaUsp|—(ni We note here that for the case of interest i which the dust is heated by (optically thin) FIR radiation. Q,=Q,.. Ma=,=vy, aud one fiuds that C~1.68 is a constant slehtly dependent ou the power law iudex of the dust opacity but otherwise independent of dust characteristics."," We note here that for the case of interest in which the dust is heated by (optically thin) FIR radiation, $Q_a=Q_e$, $\nu_a=\nu_p=\nu_e$ and one finds that $C\simeq 1.68$ is a constant slightly dependent on the power law index of the dust opacity but otherwise independent of dust characteristics."
+" For rz,22l1 the integrand in B; can be expanded as a Taylor series for small f. obtaining the approximate result Thus. the dust temperature at the center is given. to lowest order. bv heDEquisqd(XjeT1 where Ίο)Ee"," For $\tau_a\gg 1$ the integrand in ${\cal B}_i$ can be expanded as a Taylor series for small $t$, obtaining the approximate result Thus, the dust temperature at the center is given, to lowest order, by _i }, where )^4 }."
+ Notice that the depeudence of the dust temperature on the iuteusitv of the external radiation is rather weak. TyXWhe TO) ," Notice that the dependence of the dust temperature on the intensity of the external radiation is rather weak, $T_d\propto
+W^{1/(\alpha_e+4)}$ ."
+Aeain. this expression simplifies considerably when the dust is heated by (optically thick) FIR radiation. and D~0.59 in this case.," Again, this expression simplifies considerably when the dust is heated by (optically thick) FIR radiation, and $D\simeq 0.59$ in this case."
+ The above discussion has neglected the contribution of the MIR radiation., The above discussion has neglected the contribution of the MIR radiation.
+" We approximate the external MIR field with a power law spectrum in the wavelength rauge 10100 aand with a lower frequeucy cutoff at10072. with the values of 1. A, aud p given in Table DI.."," We approximate the external MIR field with a power law spectrum in the wavelength range 10–100 and with a lower frequency cut–off at, )^p, with the values of $W$, $\lambda_p$ and $p$ given in Table \ref{isf_param}."
+" Notice that p«0 in this case aud thus for the ""MIR radiation. thereis approximate cancellation iu the inteeral ou the right hand side of equation 1. between the frequency dependence of Q, and that of the external NIR field."," Notice that $p<0$ in this case and thus for the “MIR” radiation, thereis approximate cancellation in the integral on the right hand side of equation \ref{therm_eq} between the frequency dependence of $Q_{\nu}$ and that of the external MIR field."
+ The consequence is that the main contribution to the integral comes from frequencies where the optical depth is of order unity (i.e. at a waveleugth of roughly for Ay =30 mag)., The consequence is that the main contribution to the integral comes from frequencies where the optical depth is of order unity (i.e. at a wavelength of roughly for $A_{V}$ =30 mag.).
+ Substituting the above expression for the MIRfield iu eq. (1)).," Substituting the above expression for the MIRfield in eq. \ref{therm_eq}) ),"
+" we obtainwhere s—(o,|pUfa,=0.5 for our parameters. 7,) is the incomplete eauuna fuuction."," we obtainwhere $s=(\alpha_a+p+1)/\alpha_a\, = \, 0.5$ for our parameters, and $\Gamma(s,\tau_a)$ is the incomplete gamma function."
+" Here z, is s optical depth at 100 pau... which becomes ΙΤ; at lec LLL."," Here $\tau_a$ is the cloud's optical depth at 100 , which becomes unity at $A_V\simeq 414$ ."
+ Below this value. the incomplete eamunua function can be expauded for small τι. obtaining," Below this value, the incomplete gamma function can be expanded for small $\tau_a$ , obtaining ."
+"Assuming [22023 to be a normal population I star. we may adopt the simple approximation given in Lang (1980) for the velocity structure of the Galactic disk which expresses the heliocentric radial velocity as a function of distance and Galactic coordinates of the star: va 2 —0v4 + AD Sin 21 Cos*b [kms]. where. νι is the heliocentric radial velocity of 122023: δν. 1s the correction term to convert from heliocentric radial velocity to velocity in the Local Standard of Rest (vj,) and is 13.19 kms! for 122023: Oort's constant. A 15 km s! kpe!: D is the distance to the star inkpe and the Galactie coordinates for 122023. (lb) = (99.30: 1.96).","Assuming I22023 to be a normal population I star, we may adopt the simple approximation given in Lang (1980) for the velocity structure of the Galactic disk which expresses the heliocentric radial velocity as a function of distance and Galactic coordinates of the star: $_{\odot}$ = $-\delta$ $_{\odot}$ + A D Sin 2l $^{2}$ b $^{-1}$ ], where, $_{\odot}$ is the heliocentric radial velocity of I22023; $\delta$ $_{\odot}$ is the correction term to convert from heliocentric radial velocity to velocity in the Local Standard of Rest $_{\rm lsr}$ ) and is 13.19 $^{-1}$ for I22023; Oort's constant, A $\simeq$ 15 km $^{-1}$ $^{-1}$; D is the distance to the star inkpc and the Galactic coordinates for I22023, (l;b) = (99.30; $-$ 1.96)."
+ Using v. 2 — 148.31 km7! (Sec., Using $_{\odot}$ = $-$ 148.31 $^{-1}$ (Sec.
+ 3.4). we derive a distance of 28.2 kpe to 122023.," 3.4), we derive a distance of 28.2 kpc to I22023."
+ Such a large distance is unphysical and suggests that [22023 lies beyond the edge of our Galaxy., Such a large distance is unphysical and suggests that I22023 lies beyond the edge of our Galaxy.
+ It is therefore. unlikely that 122023 is a normal population I B giant/supergiant star.," It is therefore, unlikely that I22023 is a normal population I B giant/supergiant star."
+ On the post- AGB evolutionary tracks of Schónnberner (1987). a post- AGB star with Το = 24000K and log g = 3.0 (Sec.," On the $-$ AGB evolutionary tracks of Schönnberner (1987), a $-$ AGB star with $_{\rm eff}$ = 24000K and log g = 3.0 (Sec."
+" 3.8) has a core mass. M, = 0.565 M..."," 3.8) has a core mass, $_{\rm c}$ = 0.565 $_{\odot}$."
+" If 122023 is a post-AGB star. then using the relation between core-mass and quiescent luminosity maximum (Wood Zarro 1981) for AGB stars. we estimate Myo, = —4.30."," If I22023 is a $-$ AGB star, then using the relation between $-$ mass and quiescent luminosity maximum (Wood Zarro 1981) for AGB stars, we estimate $_{\rm bol}$ = $-$ 4.30."
+" Applying the bolometric correction (= —2.43) for a BIIII star (Lang 1992) we derive M, = -1.87.", Applying the bolometric correction (= $-$ 2.43) for a B1III star (Lang 1992) we derive $_{\rm v}$ = $-$ 1.87.
+ The observed V magnitude (my= 12.52: Hog et al., The observed V magnitude $_{\rm V}$ = 12.52; Hog et al.
+" 2000) of the star was corrected for extinction using E(B- V) = 0.95 (A, = 2.94).", 2000) of the star was corrected for extinction using $-$ V) = 0.95 $_{\rm v}$ = 2.94).
+ Distance modulus method then yields a distance of 1.95 kpe to the star., Distance modulus method then yields a distance of 1.95 kpc to the star.
+ Heliocentric radial velocities have been derived from wavelength shifts ofthe well defined absorption and emission lines (Tables 2. 3. 4 and 5).," Heliocentric radial velocities have been derived from wavelength shifts ofthe well defined absorption and emission lines (Tables 2, 3, 4 and 5)."
+ The mean heliocentric radial velocities from the absorption and emission lines (Tables 2 and 3) are — 148.31 + 0.60 kms! and — 144.13 + 0.72 kms! respectively., The mean heliocentric radial velocities from the absorption and emission lines (Tables 2 and 3) are $-$ 148.31 $\pm$ 0.60 $^{-1}$ and $-$ 144.13 $\pm$ 0.72 $^{-1}$ respectively.
+ Radial velocity measurements of the forbidden lines have been excluded in estimating the mean., Radial velocity measurements of the forbidden lines have been excluded in estimating the mean.
+ The quoted errors refer to the probable errors of estimation., The quoted errors refer to the probable errors of estimation.
+ Fig.2 shows the radial velocity trends with respect to the equivalent widths (W 1) and lower excitation potentials (LEP) of the absorption and emission lines respectively., Fig.2 shows the radial velocity trends with respect to the equivalent widths $_{\lambda}$ ) and lower excitation potentials (LEP) of the absorption and emission lines respectively.
+ The mean heliocentric radial velocity of the [N II]. |O I] and [Fe II] lines is 2152.90 + 0.96 kms7!. [," The mean heliocentric radial velocity of the [N II], [O I] and [Fe II] lines is $-$ 152.90 $\pm$ 0.96 $^{-1}$ . ["
+S II] and lines in 122023 have a markedly different heliocentric velocity corresponding to 2171.93 + 1.36 kms!.,S II] and lines in I22023 have a markedly different heliocentric velocity corresponding to $-$ 171.93 $\pm$ 1.36 $^{-1}$ .
+ The different radial velocities argue for a nebula., The different radial velocities argue for a non-spherical nebula.
+"Additionally, we consider Low our distance estimate depends on our asstuptions about the mass and reddening.","Additionally, we consider how our distance estimate depends on our assumptions about the mass and reddening."
+ To do this. we can consider our original four paralecter likelihood grid for cach of our 59 white dwarts.," To do this, we can consider our original four parameter likelihood grid for each of our 59 white dwarfs."
+ After mareinalising out the effective temperature for cach WD individually and combining the Lkchhood of all 59 WDs in Eqn. 6..," After marginalising out the effective temperature for each WD individually and combining the likelihood of all 59 WDs in Eqn. \ref{eqn:likelihood3},"
+ we are left with the likelihood values on a grid over reddening. mass. aud distance modulus.," we are left with the likelihood values on a grid over reddening, mass, and distance modulus."
+ For cach combination of reddening aud mass on the exid. there is a corresponding one-dimensional distauce modulus Likelihood distribution.," For each combination of reddening and mass on the grid, there is a corresponding one-dimensional distance modulus likelihood distribution."
+ The most likely distauce modulis can then be found for cach mass and reddening pair., The most likely distance modulus can then be found for each mass and reddening pair.
+ The result is a surface of most Likely distance modulus values that is well approximately bv a plane. described as where the mass is in solar masses.," The result is a surface of most likely distance modulus values that is well approximately by a plane, described as where the mass is in solar masses."
+ This approximation to our likelihood. erid is accurate to 0.003 mag and is valid only within the rauecs of mass (0.500.56 AD.) and reddening (0/— 0.08)., This approximation to our likelihood grid is accurate to 0.003 mag and is valid only within the ranges of mass $0.50-0.56$ $_{\odot}$ ) and reddening $0-0.08$ ).
+ Tere we test the assumption that the WDs iu our sample are DA atmosphere WDs by fitting our sample of 59 WDs instead with DB atinosphere models (Bergeronetal. 2OLL)., Here we test the assumption that the WDs in our sample are DA atmosphere WDs by fitting our sample of 59 WDs instead with DB atmosphere models \citep{bergeron11}.
+. The SEDs at with the DB models return teiiperatures that are on average 2200 IX cooler than the fits with the DA models., The SEDs fit with the DB models return temperatures that are on average $2200$ K cooler than the fits with the DA models.
+ These cooler temperatures result in objects that are fainter and thus our true distance modulus is closer at a value of 12.90., These cooler temperatures result in objects that are fainter and thus our true distance modulus is closer at a value of $12.90$.
+ If we compare the likchhoods of these fits. 53 of our 59 WDs have a higher likelihood value for our best fit DA moclel. with 18 being better by a factor of LO or more aud 22 being better by a factor of 10+ or more.," If we compare the likelihoods of these fits, 53 of our 59 WDs have a higher likelihood value for our best fit DA model, with 48 being better by a factor of 10 or more and 22 being better by a factor of $10^4$ or more."
+ The 6 remaining WDs that were fit better by a DB model have au average likehhood that is better by a factor of2., The 6 remaining WDs that were fit better by a DB model have an average likelihood that is better by a factor of 2.
+ Additionally. the standard deviation of the differeuces between the best fit model and all 59 WDs is 0.11 mae for the DA models and 0.18 mag for the DB models.," Additionally, the standard deviation of the differences between the best fit model and all 59 WDs is $0.11$ mag for the DA models and $0.18$ mag for the DB models."
+ Iu the field population of WDs. Tremblay&Bergeron(2008) found the ratio of DB to DA to be approximately 25% in the teuperature range of our WD sample.," In the field population of WDs, \cite{tremblay08} found the ratio of DB to DA to be approximately $25\%$ in the temperature range of our WD sample."
+ We he effect on the determined distance modulus if © of our suuple were instead DB WDs. bv randoily fittiug 25% of our WDs by DD inodels and the rest with DA models.," We test the effect on the determined distance modulus if $25\%$ of our sample were instead DB WDs, by randomly fitting $25\%$ of our WDs by DB models and the rest with DA models."
+ Uuder these conditions. the distance modulus is 13.32 with a random uncertainty of 0.06. casily cucompassing our distance modulus assuming all DÀ WDs.," Under these conditions, the distance modulus is $13.32$ with a random uncertainty of $0.06$, easily encompassing our distance modulus assuming all DA WDs."
+ The statistical error of 0.02 that is quoted from our ΠΠη likelihood analysis in Section 5.1 are lo upper and lower errors for the distance modu]nus calculated directly from the sumuned likelihood. «istribution iu Fie. r.., The statistical error of $0.02$ that is quoted from our maximum likelihood analysis in Section \ref{sec:likelihood} are $1\sigma$ upper and lower errors for the distance modulus calculated directly from the summed likelihood distribution in Fig. \ref{fig:dm}.
+" Tje systematic errors were deterined eutirely from the ""ucertainties iu our photometric calibration taken frouo Waliraioetal.(2001).", The systematic errors were determined entirely from the uncertainties in our photometric calibration taken from \cite{kalirai11}.
+. We generate 10 svuthetic WD SEDs from the models o Tremblayetal.(2011) eveuly spaced between S00)32000. EK. πρΙΟ) he raugeanee of effective teiiperatiteniperre of our WD siuuple.," We generate 10 synthetic WD SEDs from the models of \cite{tremblay11} evenly spaced between $8000-32000$ K, spanning the range of effective temperature of our WD sample."
+ We performed the same Likelihood analysis as im Section 5.1 using the uucertaiutics in the plotometric calibration as the errors on cach magnitude., We performed the same likelihood analysis as in Section \ref{sec:likelihood} using the uncertainties in the photometric calibration as the errors on each magnitude.
+ For cach of the 10 WDs. we extract a distance modulus frou their likelihood distributions along with the width of the distribution as the eror.," For each of the 10 WDs, we extract a distance modulus from their likelihood distributions along with the width of the distribution as the error."
+ The widths of the likelihood distributions ασ be due solely to our plotometric calibration uncertainties. as these svuthetic objects come directly from the model grids.," The widths of the likelihood distributions must be due solely to our photometric calibration uncertainties, as these synthetic objects come directly from the model grids."
+ Figure δ shows the systematic errors as a function of effective temperature for the svuthetic WDs ft with a cubic spline function., Figure \ref{fig:errors} shows the systematic errors as a function of effective temperature for the synthetic WDs fit with a cubic spline function.
+ For cach of our 59 WDs. we used the effective teniperature to obtain the systematic error on its determined distance modulus from the spline ft in Fie. &.," For each of our 59 WDs, we used the effective temperature to obtain the systematic error on its determined distance modulus from the spline fit in Fig. \ref{fig:errors}."
+.For the final calibration error ou the distance modulus of 17 Tuc. we quote the mean svstematic error of the 59 WDs.," .For the final calibration error on the distance modulus of 47 Tuc, we quote the mean systematic error of the 59 WDs."
+ Using this technique. we have etermiued the true distance modulus to LF Tue to he 13.36+ 0.02(randomjto.0G6(svstematic) corresponding to a distance of L70+0.0 οτι). (systematic) κρο," Using this technique, we have determined the true distance modulus to 47 Tuc to be $13.36 \pm0.02$ $ \pm 0.06$ (systematic) corresponding to a distance of $4.70\pm0.04$ $\pm0.13$ (systematic) kpc."
+ An unknown fraction of DD objects in our sample will svstematically change this result. eiviug us a smaller distance modulus. aud wider error bars as the fraction increases.," An unknown fraction of DB objects in our sample will systematically change this result, giving us a smaller distance modulus, and wider error bars as the fraction increases."
+ The closer distauce results from the fact that DB white dwarts of roughly the same SED shape are brighter than DAs. meaning we would have to fit them closer iu order to match the observed magnitudes.," The closer distance results from the fact that DB white dwarfs of roughly the same SED shape are brighter than DAs, meaning we would have to fit them closer in order to match the observed magnitudes."
+ The wider error bars result from the fact that DD models alinost universally fit our objects worse than DA models., The wider error bars result from the fact that DB models almost universally fit our objects worse than DA models.
+ Additionally. we dout know which objects arc DBs. and so this causes a systematic error that ueeds to be addressed by iterating our fitting routine while randomlyassigning some fraction of our objects to be fit as DBs rather than DAs each time.," Additionally, we don't know which objects are DBs, and so this causes a systematic error that needs to be addressed by iterating our fitting routine while randomlyassigning some fraction of our objects to be fit as DBs rather than DAs each time."
+ The fit distance modulus aud errors both change roughly linearly with DB fraction. as shown in Figure 9..," The fit distance modulus and errors both change roughly linearly with DB fraction, as shown in Figure \ref{fig:dbfraction}."
+ At 0% DBs we have the quoted fit values. at 25% DBs we have 13.32+0.06+40.06.," At $0\%$ DBs we have the quoted fit values, at $25\%$ DBs we have $13.32\pm0.06\pm0.06$."
+ We have tried a nmunber of variations in order to assess how the distance modulus chauges by varving sole of our input asstunptions., We have tried a number of variations in order to assess how the distance modulus changes by varying some of our input assumptions.
+ The whitedsvarfs forming in globular clusters today should lave masses between O510.55 M... (Reuziui&FusiPecci1988:Reuzini 1996)...," The white dwarfs forming in globular clusters today should have masses between $0.51-0.55$ $_{\odot}$ \citep{renzini88,renzini96}. ."
+ As discussed in Section 5.1... we lave used a," As discussed in Section \ref{sec:likelihood}, , we have used a"
+Tt is now established that the observed neutrinos have tiny masses and they mis with cach other [1]..,It is now established that the observed neutrinos have tiny masses and they mix with each other \cite{PDG}.
+ Theoretically. it is largely anticipated that the neutrinos are Majorana particles.," Theoretically, it is largely anticipated that the neutrinos are Majorana particles."
+ Experimental evidence for 072.) decay would deliver a conchisive confirmation of the Majorana nature of neutrinos. establishing he existence of physics bevoud the standard model (SM) [2]..," Experimental evidence for $0\nu2\beta$ decay would deliver a conclusive confirmation of the Majorana nature of neutrinos, establishing the existence of physics beyond the standard model (SM) \cite{Vogel:2006sq}."
+ An extended version of the SM. could coutain tiny nourcnormalizable terms that violate epton uuuber (LN) aud allow the 072.) decay., An extended version of the SM could contain tiny nonrenormalizable terms that violate lepton number (LN) and allow the $0\nu2\beta$ decay.
+" Probable mechanisiuus of LN violation may include exchanges by: Majorana neutrinos vays —[3.l] (the referred iiechauisii after the observation of ucutrino oscillations |1])). SUSY yarticles ο,δω, scalar bilinears (SBs) [τν e.g. doubly charged dileptous (the component of the ο(ο) triplet Tigges etc.)."," Probable mechanisms of LN violation may include exchanges by: Majorana neutrinos $\nu_M$ s \cite{ZKS,Doi} (the preferred mechanism after the observation of neutrino oscillations \cite{PDG}) ), SUSY particles \cite{SUSY1,SUSY}, scalar bilinears (SBs) \cite{BL}, , e.g. doubly charged dileptons (the component of the $SU(2)_L$ triplet Higgs etc.),"
+ leptoquarks (LQs) [8].. Hg bosous |L.9] ete.," leptoquarks (LQs) \cite{LQ}, right-handed $W_R$ bosons \cite{Doi,HKP} etc."
+ From these particles light vs are much lighter han the electron aud others are much heavier than the proton that gives two xossible classes of mechiauisius for the 072. decay: long range (vith the liebt vs in the intermediate state) and short range mechanism., From these particles light $\nu$ s are much lighter than the electron and others are much heavier than the proton that gives two possible classes of mechanisms for the $0\nu2\beta$ decay: long range (with the light $\nu$ s in the intermediate state) and short range mechanism.
+ Our adn was to exanune the possibility to discriminate among the various possible niechauisuis contributing to the 0v2-decars aud the various sources of LN violation using the information on the augular correlation of the final clectrous., Our aim was to examine the possibility to discriminate among the various possible mechanisms contributing to the $0\nu2\beta$ -decays and the various sources of LN violation using the information on the angular correlation of the final electrons.
+ We published a preliminary study aloug these lines in Ref., We published a preliminary study along these lines in Ref.
+ [10]. aud a more detailed study in Ref. |l1].., \cite{Ali:2006iu} and a more detailed study in Ref. \cite{ABZ_PRD}.
+ ere. we sumunarize the main results of Ref. |11]..," Here, we summarize the main results of Ref. \cite{ABZ_PRD}."
+" For the decay mediated by light ways. the most general effective Laerangian is the Lorentz invariant combination of the leptouic j,, aud the hadronic οι currents of definite tensor structure and chirality [12.13] where the hadronic and leptonic currents are defined as: J!=vO, and J= Onn; the leptouic currents coutain neutrino mass clgcustates. and the index / ruus over the helt cigcustates."," For the decay mediated by light $\nu_M$ s, the most general effective Lagrangian is the Lorentz invariant combination of the leptonic $j_\alpha$ and the hadronic $J_\alpha$ currents of definite tensor structure and chirality \cite{Limits,Gamov}
+ where the hadronic and leptonic currents are defined as: $J^+_\alpha=\bar{u} O_\alpha d$ and $j^i_\beta =\bar{e} O_\beta
+\nu_i$ ; the leptonic currents contain neutrino mass eigenstates, and the index $i$ runs over the light eigenstates."
+" Tere and thereafter. à sununatiou over the repeated iudices is assimucd: o. J—VxA. SEP. Trp (Or,=20!D, ofi""= iU P,=AFτα) is the "," Here and thereafter, a summation over the repeated indices is assumed; $\alpha$ $\beta$ $V\!\mp\!A$ $S\!\mp\!P$ $T_{L,R}$ $O_{T_\rho}=2\sigma^{\mu\nu}P_\rho$, $\sigma^{\mu\nu}=\frac{i}{2}\left[\gamma^\mu,\gamma^\nu\right]$ , $P_\rho=(1\mp \gamma_5)/2$ is the projector, $\rho=L,\,R$ ); the prime indicates the summation over all the Lorentz invariant contributions, except for $\alpha=\beta=V-A$ , $U_{ei}$ is the PMNS mixing matrix."
+projector.," The coefficients $\epsilon_{\alpha i}^\beta$ encode new physics, parametrizing deviations of the Lagrangian from the standard $V-A$ current-current form and mixing of the non-SM neutrinos."
+ p=L., The nonzero $\epsilon_\alpha^\beta$ for the particular SM extensions are collected in Table 1.
+ Ry: the prime indic, We have calculated the leading order in the Fermi constant and the leading contribution of the parameters $\epsilon_\alpha^\beta$ using the approximation of the relativistic electrons and nonrelativistic nucleons.
+ates — o. PERS , We take into account the $S_{1/2}$ and the $P_{1/2}$ waves for the outgoing electrons and include the finite de Broglie wave length correction for the $S_{1/2}$ wave.
+T N," Taking into account the nucleon recoil terms includingthe terms due to the pseudoscalar form factor we obtain the differential width in $\cos \theta$ for the $0^+\!(A,Z)\rightarrow\!0^+\!(A,Z+2) e^- e^-$ transitions: where $\theta$ isthe angle between the electron momenta in the rest frame of the parent nucleus, $M_{GT}$ is the Gamow–Teller nuclear matrix element and the angularcorrelation coefficient is"
+The minimum luminositv of the target photon field in the UV that is needed to achieve IC’ losses similar to curvature radiation losses ((3.2)) is about the same as the observed UV luminosity.,The minimum luminosity of the target photon field in the UV that is needed to achieve IC losses similar to curvature radiation losses \ref{Lcrit}) ) is about the same as the observed UV luminosity.
+ The upscattering of soft N-rax photons. even though the X-ray flux is hieher than the UV fIux. does not contribute much to the radiative loss of the primary. beam (XN suppression) because the observed X-ray Iuminositv is five orders of magnitude below the critical luminosity.," The upscattering of soft X-ray photons, even though the X-ray flux is higher than the UV flux, does not contribute much to the radiative loss of the primary beam (KN suppression) because the observed X-ray luminosity is five orders of magnitude below the critical luminosity."
+ The conclusion that the IC upseattering of UV photons and curvature radiation contribute about equal to the total loss of the primary beam means (hat both processes also contribute equally to the emitted power in the gamma-ray band., The conclusion that the IC upscattering of UV photons and curvature radiation contribute about equal to the total loss of the primary beam means that both processes also contribute equally to the emitted power in the gamma-ray band.
+ However. the (wo processes produce very different spectral features.," However, the two processes produce very different spectral features."
+ As we have shown before. curvature radiation photons can onlv be emitted with energies up to a few GeV for reasonable electric fields ancl curvature radii. seeEq. (5)).," As we have shown before, curvature radiation photons can only be emitted with energies up to a few GeV for reasonable electric fields and curvature radii, seeEq. \ref{1}) )."
+ The spectrum of the IC upscattered photons. on the other hand. extends io much higher energies.," The spectrum of the IC upscattered photons, on the other hand, extends to much higher energies."
+ This can be shown by assuming again that curvature radiation aud 1C losses are about equal. in which case the maximum Lorentz factor can still be estimated with Eq. (5)).," This can be shown by assuming again that curvature radiation and IC losses are about equal, in which case the maximum Lorentz factor can still be estimated with Eq. \ref{1}) )."
+" The maximum energv of the upscattered photons. ει, is then given by the maxinmn electron energy: 2gscale = 15QU While the maximum photon energy produced by IC scattering depends on the maximum electron energy. the total power emitted bv IC scattering is independent of the electron energv."," The maximum energy of the upscattered photons, $\epsilon_\gamma$, is then given by the maximum electron energy: _b m_e c^2 = ( 3 = 15 While the maximum photon energy produced by IC scattering depends on the maximum electron energy, the total power emitted by IC scattering is independent of the electron energy."
+ Instead the total power is determined by (he low-enerev target photons field., Instead the total power is determined by the low-energy target photons field.
+ Due to the steeply [alling IC cross-section in (he INN regime with increasing energyof the target photons xe 7. the maxinnun IC power Ly might not be determined by the peak luminosity in the spectral energy distribution of the target photons but be at lower energies: — ( suFU ampe Wi," Due to the steeply falling IC cross-section in the KN regime with increasing energyof the target photons $\propto \epsilon^{-2}$ , the maximum IC power $L_{KN}$ might not be determined by the peak luminosity in the spectral energy distribution of the target photons but be at lower energies: = ( )^2 _Tc _G"
+ Due to the steeply [alling IC cross-section in (he INN regime with increasing energyof the target photons xe 7. the maxinnun IC power Ly might not be determined by the peak luminosity in the spectral energy distribution of the target photons but be at lower energies: — ( suFU ampe Wig," Due to the steeply falling IC cross-section in the KN regime with increasing energyof the target photons $\propto \epsilon^{-2}$ , the maximum IC power $L_{KN}$ might not be determined by the peak luminosity in the spectral energy distribution of the target photons but be at lower energies: = ( )^2 _Tc _G"
+ Due to the steeply [alling IC cross-section in (he INN regime with increasing energyof the target photons xe 7. the maxinnun IC power Ly might not be determined by the peak luminosity in the spectral energy distribution of the target photons but be at lower energies: — ( suFU ampe Wig.," Due to the steeply falling IC cross-section in the KN regime with increasing energyof the target photons $\propto \epsilon^{-2}$ , the maximum IC power $L_{KN}$ might not be determined by the peak luminosity in the spectral energy distribution of the target photons but be at lower energies: = ( )^2 _Tc _G"
+LFs aro a convenient wav fo describe the galaxy population aud to eet hints about the nechauisus of formation and evolution of galaxies.,LFs are a convenient way to describe the galaxy population and to get hints about the mechanisms of formation and evolution of galaxies.
+ Two scenarios are in competition to explain the history of galaxies up to the preseut epoch., Two scenarios are in competition to explain the history of galaxies up to the present epoch.
+ The formation of cllipical galaxies is especially intriguing. because despite their old and apparently simple stellar population. their process of formation is far from being understood.," The formation of elliptical galaxies is especially intriguing, because despite their old and apparently simple stellar population, their process of formation is far from being understood."
+ The two uodoels in competition are:, The two models in competition are:
+with the period of the external shear.,with the period of the external shear.
+ However. their perturbative calculation is applicable onlv as long as the forcing periodFT is small compared to the turnover tme of (he largest eddies.," However, their perturbative calculation is applicable only as long as the forcing period$T$ is small compared to the turnover time of the largest eddies."
+ In particular the perturbative treatment is not able to provide (he maxinunm value the effective viscosity reaches and the frequency. al which it reaches it. which is of great importance in caleulating tidal interactions and dissipation of pulsations.," In particular the perturbative treatment is not able to provide the maximum value the effective viscosity reaches and the frequency at which it reaches it, which is of great importance in calculating tidal interactions and dissipation of pulsations."
+ In this article we use the Penevetal.(2008a) spectral anelastic code to perform a direct calculation of the turbulent dissipation in a convective zone. by introducing external shear as an extra body force in the fluid equations.," In this article we use the \citet{Penev_Barranco_Sasselov_08a} spectral anelastic code to perform a direct calculation of the turbulent dissipation in a convective zone, by introducing external shear as an extra body force in the fluid equations."
+ The goal is to investigate the applicability of effective viscosity as an approximation to the actual turbulent dissipation. and to derive directly. an effective viscosity. prescription and compare it against the (1997) formalism.," The goal is to investigate the applicability of effective viscosity as an approximation to the actual turbulent dissipation, and to derive directly an effective viscosity prescription and compare it against the \citet{Goodman_Oh_97} formalism."
+ The details of the numerical simulation aud (he ecuations evolved are presented in Penevelal.(2008a)., The details of the numerical simulation and the equations evolved are presented in \citet{Penev_Barranco_Sasselov_08a}.
+. We are simulating a rectangular box with impenetrable. constant temperature top and bottom boundaries (the z velocity. vanishes and (he temperature is held at some constant. value at the top and bottom walls of the box) using the anelastic approximation.," We are simulating a rectangular box with impenetrable, constant temperature top and bottom boundaries (the $\hat{z}$ velocity vanishes and the temperature is held at some constant value at the top and bottom walls of the box) using the anelastic approximation."
+ The background state. and the parameters with which all the runs presented in (his paper were computed. are (he same as (he parameters used for all couvectively unstable (ests of section 4.1 of Penevetal. (2008a)..," The background state, and the parameters with which all the runs presented in this paper were computed, are the same as the parameters used for all convectively unstable tests of section 4.1 of \citet{Penev_Barranco_Sasselov_08a}. ."
+ For convenience we remind them here:, For convenience we remind them here:
+in van Ojik a£) and the kinetic temperature γρ=107 Ix acloptecl by van OjikaL. we can estimate the range for the velocity dispersion of bulk motions: 24 km lEm<55 km 1.,"in van Ojik ) and the kinetic temperature $T_{kin} = 10^4$ K adopted by van Ojik, we can estimate the range for the velocity dispersion of bulk motions: 24 km $^{-1} \la \sigma_t \la 55$ km $^{-1}$."
+" Ht follows that afer,21: within the LIE absorption lino gas embedded: in the galactic halos.", It follows that $\sigma_t/v_{th} > 1$ within the HI absorption line gas embedded in the galactic halos.
+" Both lindings high absolute value of a and high afey), ratio. make the mesoturbulent approach to appear to be more realistic one."," Both findings -- high absolute value of $\sigma_t$ and high $\sigma_t/v_{th}$ ratio, – make the mesoturbulent approach to appear to be more realistic one."
+" Besides. our present IGMC calculations vielded for the ns=2.504 absorption svstem a,c25 km | which lies just within the observed range. whereas σι2 kms 1 found by DAFT is evidently too low."," Besides, our present RMC calculations yielded for the $z_a = 2.504$ absorption system $\sigma_t \simeq 25$ km $^{-1}$ which lies just within the observed range, whereas $\sigma_t \simeq 2$ km $^{-1}$ found by T is evidently too low."
+ Vhe main conclusion of our work is that the mesoturbulent approach together with the RAC computational scheme allows to solve the inverse. problem for the ILID Lye absorption and to restore the physica parameters of the absorbing gas as well as the projection of the velocity Geld distribution., The main conclusion of our work is that the mesoturbulent approach together with the RMC computational scheme allows to solve the inverse problem for the H+D $\alpha$ absorption and to restore the physical parameters of the absorbing gas as well as the projection of the velocity field distribution.
+ The proposed. computational procedure enables us to craw confidence regions for ciffercnt pairs of the acloptec parameters., The proposed computational procedure enables us to draw confidence regions for different pairs of the adopted parameters.
+ From the analysis of the svnthetic LL|D Lye spectrum. it was found that. in general. the measured. D/L and Nyy values are anti-correlated.," From the analysis of the synthetic H+D $\alpha$ spectrum it was found that, in general, the measured D/H and $N_{\rm HI}$ values are anti-correlated."
+ The study of the template LL]D Lye profile (which reproduces the original spectrum of BACLT) vields Ο/Η = (3.75c0.85)«107 (24)., The study of the template H+D $\alpha$ profile (which reproduces the original spectrum of T) yields D/H = $(3.75 \pm 0.85)\times10^{-5}$ $(2\sigma)$.
+ We therefore adopt the value of 4.60' às à conservative upper limit on the primordial abundance of D relative to hydrogen., We therefore adopt the value of $4.6\times10^{-5}$ as a conservative upper limit on the primordial abundance of D relative to hydrogen.
+" We conclude that the discordance of D/II with the {Πο and ""Li primordial abundances noted by D&TT is a consequence of the use of the miceroturbulent model.", We conclude that the discordance of D/H with the $^4$ He and $^7$ Li primordial abundances noted by T is a consequence of the use of the microturbulent model.
+ Within the framework of the generalized. model one finds. good agreement between the measurements mentioned above and the SBBN predictions., Within the framework of the generalized model one finds good agreement between the measurements mentioned above and the SBBN predictions.
+ This work was supported. by the Deutsche Forschunesgemeinschalt. anc by the REDIU grant No.," This work was supported by the Deutsche Forschungsgemeinschaft, and by the RFBR grant No."
+ 96-02-16905-a. The authors thank Dr. 1. Aeafonova for valuable suggestions on the RALC technique and. for her kind help in the development of the RAIC computer code., 96-02-16905-a. The authors thank Dr. I. Agafonova for valuable suggestions on the RMC technique and for her kind help in the development of the RMC computer code.
+ SAL gratefully acknowledges the hospitality of the Institut [für Theoretische. Physik der. Universititt Frankfurt am Alain., SAL gratefully acknowledges the hospitality of the Institut fürr Theoretische Physik der Universitätt Frankfurt am Main.
+ , 
+The formation of super-km-sized planetesimals is an important step towards terrestrial planets and the solid cores of gas and ice giants (e.g.Safronov.1969:Goldreichetal..2004;Chi-ang&Youdin. 2010).,"The formation of super-km-sized planetesimals is an important step towards terrestrial planets and the solid cores of gas and ice giants \citep[e.g.][]{Safronov1969,Goldreich+etal2004,ChiangYoudin2010}."
+. The asteroid and Kuiper belts of the solar system. as well as the extrasolar debris dises. are believed to be left-over populations of planetesimals that did not grow to planets.," The asteroid and Kuiper belts of the solar system, as well as the extrasolar debris discs, are believed to be left-over populations of planetesimals that did not grow to planets."
+ Comparing models and simulations of planetesimmal formation to observations of such planetesimal belts constrains our theoretical picture of the planetesimal formation stage. and at the same time 1t gives insight into the physical processes that shaped the architectures of these systems (Morbidelltietαἱ.&Trujillo.2010:Krivov.KenyonBromley. 2010).," Comparing models and simulations of planetesimal formation to observations of such planetesimal belts constrains our theoretical picture of the planetesimal formation stage, and at the same time it gives insight into the physical processes that shaped the architectures of these systems \citep{Morbidelli+etal2009,Weidenschilling2010,Nesvorny+etal2010,SheppardTrujillo2010,Krivov2010,KenyonBromley2010}."
+. Planetesimal formation takes place in a complex environment of turbulent gas interacting via drag forces with particles of many sizes., Planetesimal formation takes place in a complex environment of turbulent gas interacting via drag forces with particles of many sizes.
+ The streaming instability thrives in the systematic relative motion of gas and particles and leads to spontaneous clumping of particles (Youdin&Goodman.2005:Johansen&Youdin.2007:BaiStone. 2010b).. seeding a gravitational collapse into bound clumps etal..2009) and further to solid planetesimals (Nesvornyetal. 2010).," The streaming instability thrives in the systematic relative motion of gas and particles and leads to spontaneous clumping of particles \citep{YoudinGoodman2005,JohansenYoudin2007,BaiStone2010b}, seeding a gravitational collapse into bound clumps \citep{Johansen+etal2009} and further to solid planetesimals \citep{Nesvorny+etal2010}."
+. While the latest years have seen major progress in numerical modelling of drag force interaction betwee particles and gas (Youdin&Johansen.2007;Balsaraetal..2009;Miniati.2010:Bai&Stone.2010a) às well as the self-gravity of the particle layer (Johansenetal..2007;Reietal. 2010).. good algorithms for treating simultaneously hydrodynamies. gravitational dynamies and particle collisions are still missing.," While the latest years have seen major progress in numerical modelling of drag force interaction between particles and gas \citep{YoudinJohansen2007,Balsara+etal2009,Miniati2010,BaiStone2010a} as well as the self-gravity of the particle layer \citep{Johansen+etal2007,Rein+etal2010}, good algorithms for treating simultaneously hydrodynamics, gravitational dynamics and particle collisions are still missing."
+ There are two main approaches in astrophysics to treating particle collisions in numerical simulations., There are two main approaches in astrophysics to treating particle collisions in numerical simulations.
+ Modelling a set of with collision tracking allows simulation of particle aggregation m close concordance with the nature of real physical collisions., Modelling a set of with collision tracking allows simulation of particle aggregation in close concordance with the nature of real physical collisions.
+ This method has successfully been applied to model the particle rings of Saturn (Wisdom&Tremaine.1988:Salo.1991:Karjalainen&2004) and to model collisions between individual dust grains and aggregates (Dominik&Nübold.2002).," This method has successfully been applied to model the particle rings of Saturn \citep{WisdomTremaine1988,Salo1991,KarjalainenSalo2004} and to model collisions between individual dust grains and aggregates \citep{DominikNubold2002}."
+ The drawback of the physical-particle approach is that the size of the system is limited by the number of numerical particles that can be afforded in the simulation., The drawback of the physical-particle approach is that the size of the system is limited by the number of numerical particles that can be afforded in the simulation.
+ The formation of a Ceres-mass planetesimal from 10-em-sized rocks would rrequire tracking of O(107) particles. orders of magnitude beyond what current computational resources allow.," The formation of a Ceres-mass planetesimal from 10-cm-sized rocks would require tracking of $\mathcal{O}(10^{20})$ particles, orders of magnitude beyond what current computational resources allow."
+ Algorithms involving group collections of physical particles into much larger numerical particles under conservation of total mass M and mean free path 1., Algorithms involving group collections of physical particles into much larger numerical particles under conservation of total mass $M$ and mean free path $\lambda$.
+ Decreasing the particle number N to a number that can be handled in à computer simulation. while maintaining 4!=(N/V by artificially increasing the collisional cross section c. yields the correct collision frequency m systems that are much larger than what can be resolved with the physical particle approach.," Decreasing the particle number $N$ to a number that can be handled in a computer simulation, while maintaining $\lambda^{-1} \equiv (N/V) \sigma$ by artificially increasing the collisional cross section $\sigma$, yields the correct collision frequency in systems that are much larger than what can be resolved with the physical particle approach."
+ The inflated particle approach was used recently by &Chiang (2007). Michikoshietal. (2007).. Nesvornyetal. (2010).. and Reinetal.(2010).. with different methods for tracking the actual collision. but the concept of bloated particles has deeper roots (e.g.Kokubo&Ida.1996).," The inflated particle approach was used recently by \cite{LithwickChiang2007}, \cite{Michikoshi+etal2007}, \cite{Nesvorny+etal2010}, and \cite{Rein+etal2010}, with different methods for tracking the actual collision, but the concept of bloated particles has deeper roots \citep[e.g.][]{KokuboIda1996}."
+. In this paper we put forward a new algorithm to model collisions between numericalsuperparticles., In this paper we put forward a new algorithm to model collisions between numerical.
+ Superparticles are designed to represent swarms of physical particles., Superparticles are designed to represent swarms of physical particles.
+ The aerodynamical properties of the superparticle tthe friction time) is still that of a single physical particle., The aerodynamical properties of the superparticle the friction time) is still that of a single physical particle.
+ Superparticles are widely used to model the solid particle component in computer simulations of coupled gas and particle motion in protoplanetary discs (Johansen&Youdin.2007;BatStone. 2010b).," Superparticles are widely used to model the solid particle component in computer simulations of coupled gas and particle motion in protoplanetary discs \citep{JohansenYoudin2007,BaiStone2010b}."
+. Since superparticles can be considered to represent swarms of smaller particles. direct collision tracking is not possible.," Since superparticles can be considered to represent swarms of smaller particles, direct collision tracking is not possible."
+ Johansenetal.(2007) modelled superparticle collisions by damping the random motion of particles inside a grid cell on the collisional time-scale., \cite{Johansen+etal2007} modelled superparticle collisions by damping the random motion of particles inside a grid cell on the collisional time-scale.
+ They showed that inelastic collisions. where part of the kinetic energy 1s converted to heat and deformation duringthe collisions. is beneficial for the gravitational collapse and allows the formation of planetesimals in protoplanetary dises of lower mass. compared to simulations without damping.," They showed that inelastic collisions, where part of the kinetic energy is converted to heat and deformation duringthe collisions, is beneficial for the gravitational collapse and allows the formation of planetesimals in protoplanetary discs of lower mass, compared to simulations without damping."
+ However. the simplified collision scheme of Johansenetal.(2007) is," However, the simplified collision scheme of \cite{Johansen+etal2007} is"
+deeper iu sensitivity than the survey.,deeper in sensitivity than the survey.
+ Some 17 objects were fouud to coincide with aSiubad source., Some 17 objects were found to coincide with a source.
+" However. such a correlation could be muisleacing as most of the distances to the sources are too large (> GU"") to be considered reliable."," However, such a correlation could be misleading as most of the distances to the sources are too large $>60''$ ) to be considered reliable."
+" Onlv three sources have a distance to aSimbad source of «50"" aud mav be considered to be ideutified.", Only three sources have a distance to a source of $<50''$ and may be considered to be identified.
+ This is a very siall fraction of all catalogued sources here., This is a very small fraction of all catalogued sources here.
+ Four sources are found close to the immer ring of thePSPC window support svsteii aud may be artifacts., Four sources are found close to the inner ring of the window support system and may be artifacts.
+ A more thorough investigation of these sources appears to be required as these sources can still be real., A more thorough investigation of these sources appears to be required as these sources can still be real.
+ Another 16 sources have been found inside the iuner support ring of thePSPC detector and are considered as fiui caucliclates., Another 46 sources have been found inside the inner support ring of the detector and are considered as firm candidates.
+ This comprises 0: all reliable cuties iu the catalogue., This comprises of all reliable entries in the catalogue.
+" Sources found outside the detector ring suffer due to less accurate positions. """, Sources found outside the detector ring suffer due to less accurate positions. “
+"Sereenius of the catalog using harduess ratios derived from stunulated power-law slope (LS spectra for hard N-rav binaries. and powoer-huv slope 2.0. and slope 2.6 spectra for ""yadio loud” and ""radio quiet? ACUNS respectively gives 13 firm weak hid N-vav binary candidates.","Screening” of the catalog using hardness ratios derived from simulated power-law slope –0.8 spectra for hard X-ray binaries, and power-law slope –2.0, and slope –2.6 spectra for “radio loud” and “radio quiet” AGNs respectively gives 43 firm weak hard X-ray binary candidates."
+ Six previous hard X-ray binary candidates are consistent with ACNs and Ll caucidates are consistent with either class (ef., Six previous hard X-ray binary candidates are consistent with AGNs and 11 candidates are consistent with either class (cf.
+ Table 1)., Table 1).
+" Recent observations towards the SAIC with.ASC'A,BeppoSAX aud established 10 new. X-ray pulsus in this galaxy (Tab.1)."," Recent observations towards the SMC with, and established 10 new X-ray pulsars in this galaxy (Tab.4)."
+ Most Gf not all) of them appear to be connected with a De-tvpe donor star., Most (if not all) of them appear to be connected with a Be-type donor star.
+ Pulsation periods in the rauge 3-315 s have beeu determined., Pulsation periods in the range 3-345 s have been determined.
+ This rauge in pulsation periods is covered by the rauge of oulsatioun periods found iu the galactic Be-N-rav binarics of d 1500 8 (cf., This range in pulsation periods is covered by the range of pulsation periods found in the galactic Be-starX-ray binaries of $\sim$ 4–1500 s (cf.
+ van deu Heuvel Rappaport 1987)., van den Heuvel Rappaport 1987).
+ An orbital period has been estimated oulv for thetwo systems AN J0051-722 and XTE Jo053-τοι with 110-120 aud 139 days., An orbital period has been estimated only for the two systems AX J0051-722 and XTE J0053-724 with 110-120 and 139 days.
+ The orbital period deduced for AN J0051-722 is iu aegreciment with the relation between pulsation period aud orbital period fouud bv Corhet (1986). while the orbital period estimated for NTE J0053-721 is twice the predicted period.," The orbital period deduced for AX J0051-722 is in agreement with the relation between pulsation period and orbital period found by Corbet (1986), while the orbital period estimated for XTE J0053-724 is twice the predicted period."
+ Five of these new XNaaw pulsus Way have a counterpart in our SMC N-ray catalogue., Five of these new X-ray pulsars may have a counterpart in our SMC X-ray catalogue.
+ Source 98 (RN JO053.9-7226) has heen discovered withRATE (cf., Source 98 (RX J0053.9-7226) has been discovered with (cf.
+ Levine et al., Levine et al.
+ 1996) in an outburst aud 16.6 s pulsations have beeu found (Corbet ct al., 1996) in an outburst and 46.6 s pulsations have been found (Corbet et al.
+ 1998)., 1998).
+ This coufiruis the correct classification of this souree., This confirms the correct classification of this source.
+ In addition. source 79 (RN JOO51.3-721¢j) lav be ideutical with AN JOO51-722 (Corbet ct al.," In addition, source 79 (RX J0051.3-7216) may be identical with AX J0051-722 (Corbet et al."
+ 1998) also confirming the correct classification., 1998) also confirming the correct classification.
+ Source 89 coincides with the trausieut source of Cowley et al. (, Source 89 coincides with the transient source of Cowley et al. (
+1997) and AN J0051-73.3 (Yokogawa et al.,1997) and AX J0051-73.3 (Yokogawa et al.
+ 1998b) coincides with RX JOO5O.7-7316 (cf., 1998b) coincides with RX J0050.7-7316 (cf.
+ Cook. 1998: Cowley et al., Cook 1998; Cowley et al.
+ 1997: IKaliabka 1998)., 1997; Kahabka 1998).
+ This source with catalog iudex 72 apparently both fits to the AGN aud the lard ταν binary class but if is a strong candidate for a lard N-rav biniurv., This source with catalog index 72 apparently both fits to the AGN and the hard X-ray binary class but it is a strong candidate for a hard X-ray binary.
+ The derived umber of 51 hare X-ray binary candidates may be compared with the uumuber of X-rav binaries precicted from the population svuthesis calculations of Dalton and Sarazin (1995) for the SMC., The derived number of 51 hard X-ray binary candidates may be compared with the number of X-ray binaries predicted from the population synthesis calculations of Dalton and Sarazin (1995) for the SMC.
+ These calculations predict [6 binaries with luminosities iu excess of 10?!cres 1. the lower scusitivity lait of our SMC N-vay survey.," These calculations predict 46 binaries with luminosities in excess of $10^{34}\ {\rm erg}\ {\rm s}^{-1}$ , the lower sensitivity limit of our SMC X-ray survey."
+ 2 (e.g...27???) +—2," $z\simeq2$ \citep[e.g.,][]{Daddi2005,Trujillo2007,vanDokkum2008,Buitrago2008,Damjanov2009}
+ $z\sim2$ \citep[e.g.,][hereafter T05]{Treu2005}. \citealt{Hopkins2009};"
+ (?:: see? (o...2)..," \citealt{Cassata2009} \citep[][and references therein]{Khochfar2006,Naab2009,Hopkins2010}. \citep[e.g.,][]{Muzzin2009}."
+ out to 2~1 (TOS: ?.. hereafter vaWs). sueeestingOO a sinall but detectable difference in average size at fixed mass when compared to the local universe.," out to $z\sim1$ (T05; \citealt{vanderWel2008}, hereafter vdW08), suggesting a small but detectable difference in average size at fixed mass when compared to the local universe."
+ But bevoud Docl. there is ttle liel-quality dynamical data for field. spheroidals.," But beyond $z\simeq1$, there is little high-quality dynamical data for field spheroidals."
+" ο undertook an heroic observation of a single i>2 source with a stellar mass ~2«&102AJ, and an effective radius r,=OLS kpe typical of compact galaxies at 2τε2.3."," \citet{vanDokkum2009}
+ undertook an heroic observation of a single $z>2$ source with a stellar mass $\simeq 2 \times 10^{11}M_\odot$ and an effective radius $r_e=0.8$ kpc typical of compact galaxies at $z\simeq2.3$."
+" The spectrum has a claimed. stellar velocity: dispersion.⋅ of e=510!""EIEoo’ kis oL suggestingOO a reniwhkably deuse system."," The spectrum has a claimed stellar velocity dispersion of $\sigma=510^{+165}_{-95}$ km $^{-1}$, suggesting a remarkably dense system."
+ van Dokkumua ct al., van Dokkum et al.
+ postulate the initial dissipative ollapse at z23 of a high mass “core” but are unable to account for its subsequent evolution outo the 2z— scaliug relations., postulate the initial dissipative collapse at $z\simeq3$ of a high mass “core” but are unable to account for its subsequent evolution onto the $z\simeq1$ scaling relations.
+ The quantitative effect of ninor mergers on the plivsical size of a galaxy involves πας variables. aud it 15 unclear whether such draatic size evolution is possible while maintaining the tightucss of the fundamental plane aud its projections (?)..," The quantitative effect of minor mergers on the physical size of a galaxy involves many variables, and it is unclear whether such dramatic size evolution is possible while maintaining the tightness of the fundamental plane and its projections \citep{Nipoti2009}."
+ Tuterpretation of the observed treuds at fixed as further complicated by the so-called “progenitor bias” (2): if galaxies exc»v bv div iiergers. the main progenitor of a present-day nassive ealaxy did not have the same lnass at zo2.," Interpretation of the observed trends at fixed is further complicated by the so-called “progenitor bias” \citep{vanderWel2009}: if galaxies grow by dry mergers, the main progenitor of a present-day massive galaxy did not have the same mass at $z\sim2$."
+ Similarly. if galaxies become recognizable as spheroidals onv above a certain threshold in stellar velocity dispersio1i cpgr that depends on redshift. it is clear that the acdition of a new acl less deuse xopulation could winuc a false evolutionary treud.," Similarly, if galaxies become recognizable as spheroidals only above a certain threshold in stellar velocity dispersion $\sigma_{\rm ET}$ that depends on redshift, it is clear that the addition of a new – and less dense – population could mimic a false evolutionary trend."
+ This yas can be reduced by considering galaxy sizes at fixed c., This bias can be reduced by considering galaxy sizes at fixed $\sigma$.
+" Foremost. σ canges very little uuder a variety of erowthn mechanisius (c.e..7) and it is therefore a etter “label” thaMay, to track the assembly. history."," Foremost, $\sigma$ changes very little under a variety of growth mechanisms \citep[e.g.,][]{Hopkins2010} and it is therefore a better “label” than to track the assembly history."
+Secondly.σ is closolv correlated with stellar age (7). aud lherefore offers the most direct wayto track theevolving »opulation.,"Secondly,$\sigma$ is closely correlated with stellar age \citep{vanderWel2009} and therefore offers the most direct wayto track theevolving population."
+ Given there is uo clear conscnsus in unuderstaudiue the, Given there is no clear consensus in understanding the
+Standard: proto-planetary disc models (Chiang&Col-dreich.LOOT) show that the inner ~0.1 AU region is too hot to allow existence of small solid. particles there.,"Standard proto-planetary disc models \citep{CG97} show that the inner $\sim
+0.1$ AU region is too hot to allow existence of small solid particles there."
+ Thus planets should not be able to grow there., Thus planets should not be able to grow there.
+ Yet observations of nearby solar type stars show that many of them clo host planets in that inhospitable to planet formation region., Yet observations of nearby solar type stars show that many of them do host planets in that inhospitable to planet formation region.
+ The very first exoplanet to be convincingly detected: had. the separation of Ro0.05 AU to its parent star and. had a mass of about that of Jupiter (Mayor&Queloz1995)., The very first exoplanet to be convincingly detected had the separation of $R \sim 0.05$ AU to its parent star and had a mass of about that of Jupiter \citep{MQ95}.
+. Such gas giant planets circling their parent stars in a close oximity (A0.1 AW) are now called: shot jupiters”., Such gas giant planets circling their parent stars in a close proximity $R\simlt 0.1$ AU) are now called “hot jupiters”.
+ Lt is xdieved that they are explained by the inward. racial cirift (migration) of the planets born further out (Linetal.L996)., It is believed that they are explained by the inward radial drift (migration) of the planets born further out \citep{Lin96}.
+. TheAÁepíer mission has recently produced. a number of surprising results (Boruckictal.2011).. one of which is hat there is an even greater number of smaller. planets. cem. Earth-size to Neptune-size. in that region.," The mission has recently produced a number of surprising results \citep{BoruckiEtal11}, one of which is that there is an even greater number of smaller planets, e.g., Earth-size to Neptune-size, in that region."
+" Lt is similarly obvious that these smaller ""hot planets also aad to be brought there from further out. by an inward racial migration.", It is similarly obvious that these smaller “hot” planets also had to be brought there from further out by an inward radial migration.
+ One practical cilliculty in testing this idea. though. is that the migration. of smaller planets is expected. to occur in the poorly understood. “type D regime (Paardekooper&Papaloizou2008).," One practical difficulty in testing this idea, though, is that the migration of smaller planets is expected to occur in the poorly understood “type I” regime \citep{PP08}."
+. This is cillerent from the better. understood. “type H. regime by which the giant planets. migrate (Lin&Papaloizou1986)., This is different from the better understood “type II” regime by which the giant planets migrate \citep{LinPap86}.
+. This theoretical clilliculty leads to a large range in uncertainties in the predictions of the detailed. Core Accretion mocel calculations (c.g..seeFie.5inIda&Lin2008).," This theoretical difficulty leads to a large range in uncertainties in the predictions of the detailed Core Accretion model calculations \citep[e.g., see Fig. 5
+ in][]{IdaLin08}."
+. ltecently. the kev importance of the radial migration of the earliest. gas condensations formed in the massive dises by the gravitational instability was realised (Vorobvoyv&Basu2006:Boleyctal.2010).," Recently, the key importance of the radial migration of the earliest gas condensations formed in the massive discs by the gravitational instability was realised \citep{VB06,BoleyEtal10}."
+".. Analytical estimates (Navakshin2010a) ancl numerical simulations (Vorobvoyv&Basu2006.2010:Boleyetal.ChaNavakshin2010) show that these condensations can migrate all the way from their birth-place in the outer &~LOO AU disc into the inner  few AU cise and be disrupted there during the earliest massive dise phase (25 LO"" ves. vpically)."," Analytical estimates \citep{Nayakshin10c} and numerical simulations \citep{VB06,VB10,BoleyEtal10,ChaNayakshin10} show that these condensations can migrate all the way from their birth-place in the outer $R\sim 100$ AU disc into the inner $\sim$ few AU disc and be disrupted there during the earliest massive disc phase $t\simlt$ $\times 10^5$ yrs, typically)."
+ Lt was pointed out by Boleyetal.(2010) and Navakshin(2010a) that this migration-ancd-cdisruption sequence vields an unexplored way of forming terrestrial like planets., It was pointed out by \cite{BoleyEtal10} and \cite{Nayakshin10c} that this migration-and-disruption sequence yields an unexplored way of forming terrestrial like planets.
+ 1 dust. &rows anc sediments in the centre of the clump and orms a solid density core there. then tidal disruption of he gas clump may leave a solid core an Earth-like proto-Xeanets (notetheconnectiontoearlierideasofMeCrea&Williams1965:Boss1998:ctal. 20," If dust grows and sediments in the centre of the clump and forms a solid density core there, then tidal disruption of the gas clump may leave a solid core – an Earth-like proto-planets \citep[note the
+ connection to earlier ideas
+ of][]{McCreaWilliams65,Boss98,BossEtal02}."
+02)... Navakshin(2010c.b) used a simple spherically symmetric. radiation ivdrocyvnamic| code with the dust. grains as à second.[uid o delineate the conditions when such a mechanism for the solid core growth can work.," \cite{Nayakshin10a,Nayakshin10b}
+ used a simple spherically symmetric radiation hydrodynamic code with the dust grains as a secondfluid to delineate the conditions when such a mechanism for the solid core growth can work."
+ Based on the potential promise of these ideas. Navakshin(2010a) formulated the 7Tidal Downsizine” (PDhereafter:Navakshin2010a) hypothesis or planet formation.," Based on the potential promise of these ideas, \cite{Nayakshin10c} formulated the “Tidal Downsizing” \citep[TD
+ hereafter;][]{Nayakshin10c} hypothesis for planet formation."
+ In this picture a partial disruption of a 10M; eas clump Ovhich we also call giant embryos: GI) caves a giant planet. whereas a complete disruption vields a terrestrial like planet.," In this picture a partial disruption of a $\sim 10 M_J$ gas clump (which we also call giant embryos; GEs) leaves a giant planet, whereas a complete disruption yields a terrestrial like planet."
+ In this Letter we continue to assess the potential utility of the TD hypothesis to planet. formation., In this Letter we continue to assess the potential utility of the TD hypothesis to planet formation.
+ We note that another ingredient. muted but not explicitly. considered. by Boleyetal.(2010):Navakshin (2010c.b.a).. must be included in the scheme.," We note that another ingredient, muted but not explicitly considered by \cite{BoleyEtal10,Nayakshin10a,Nayakshin10b,Nayakshin10c}, , must be included in the scheme."
+ To explain it. consider isolated Ces first.," To explain it, consider isolated GEs first."
+"The plot shows the forced eccentricity, as a function of the ratio a‘/ag, calculated with three different methods.","The plot shows the forced eccentricity, as a function of the ratio $a^*/a_B$, calculated with three different methods."
+ Recall that this model predicts a linear increase of ey with the semimajor axis., Recall that this model predicts a linear increase of $e_f$ with the semimajor axis.
+" Finally, the value of the forced eccentricity determined with our second-order Hamiltonian F* is shown as a continuous curve."," Finally, the value of the forced eccentricity determined with our second-order Hamiltonian $F^*$ is shown as a continuous curve."
+" The agreement with the numerical data is very good, and the saturation in the value of ey is reproduced quite well."," The agreement with the numerical data is very good, and the saturation in the value of $e_f$ is reproduced quite well."
+" Since we have avoided all small denominators in the generating function B,, the model curve is smooth and shows no indication of the effects of mean-motion resonances."," Since we have avoided all small denominators in the generating function $B_1$, the model curve is smooth and shows no indication of the effects of mean-motion resonances."
+" As mentioned before, although the second-order model leads to significant improvement in the secular solution, as well as allowing the magnitude of the short-period orbital variations to be modeled, it is much too complex to constitute a workable model."," As mentioned before, although the second-order model leads to significant improvement in the secular solution, as well as allowing the magnitude of the short-period orbital variations to be modeled, it is much too complex to constitute a workable model."
+" For this reason, we wondered whether the empiric correction term introduced by Thébbault et al. ("," For this reason, we wondered whether the empiric correction term introduced by Thébbault et al. ("
+2006) for the secular frequency could be extended to reproduce both the forced eccentricity and the short-period variations.,2006) for the secular frequency could be extended to reproduce both the forced eccentricity and the short-period variations.
+" Of course it is not expected to yield the exact same results, but if the errors are not significant, such an empirical second-order approximation could constitute a simple quantitative analytical model."," Of course it is not expected to yield the exact same results, but if the errors are not significant, such an empirical second-order approximation could constitute a simple quantitative analytical model."
+ Following the same approach as Thébbault et al. (, Following the same approach as Thébbault et al. (
+"2006), we use ey, and go to denote the first-order expressions for the forcedeccentricity and secular frequency, and reserve ey and g for the second-order values.","2006), we use ${e_f}_0$ and $g_0$ to denote the first-order expressions for the forcedeccentricity and secular frequency, and reserve $e_f$ and $g$ for the second-order values."
+" The idea then is to write ey=€rg(1+dey) (and a similar equation for g), and attempt to model the correction terms óey and óg."," The idea then is to write $e_f = {e_f}_0 ( 1 + \varepsilon \delta e_f)$ (and a similar equation for $g$ ), and attempt to model the correction terms $\delta e_f$ and $\delta g$."
+" After several tests and multivariate linear regressions, we find that the expressions ey epo 42) (ῴλα- [iam g £0 (os) ( £)*a a agree with the complete second-order model very closely."," After several tests and multivariate linear regressions, we find that the expressions e_f _0 ) ( )^2 ] g g_0 ) ( )^2 ] agree with the complete second-order model very closely."
+ There are some slight differences in g with respect to the original formula introduced by Thébbault et al. (, There are some slight differences in $g$ with respect to the original formula introduced by Thébbault et al. (
+2006) but they are minor and not very significant.,2006) but they are minor and not very significant.
+" Finally, the expressions for 6/0 and go are those given in (3)) and (4))."," Finally, the expressions for ${e_f}_0$ and $g_0$ are those given in \ref{eq3}) ) and \ref{eq4}) )."
+ In terms of (17)) the secular Hamiltonian may be approximated well by Fax vate Hk? + h*?) _ eit’) where k*=e*cos(A@”*) and h*=e*sin(Aw”) are the new regular secular variables., In terms of \ref{eq17}) ) the secular Hamiltonian may be approximated well by F^* n^* ^2 g ^2 + ^2) - e_f k^* ] where $k^* = e^*\cos{(\Delta \varpi^*)}$ and $h^* = e^*\sin{(\Delta \varpi^*)}$ are the new regular secular variables.
+" Finally, the semi-amplitude of the short-period variations in eccentricity can also be empirically modeled according to the expression Δε vcn (ca) Figure 3 once again compares the estimated values of ey and g, this time for three different values of the eccentricity eg of the binary component."," Finally, the semi-amplitude of the short-period variations in eccentricity can also be empirically modeled according to the expression e 10 ( ) ( )^3 Figure \ref{fig3} once again compares the estimated values of $e_f$ and $g$ , this time for three different values of the eccentricity $e_B$ of the binary component."
+ Given the simplicity, Given the simplicity
+The optical polarimetric campaign was conducted in early September 2008. between MJD 54710-54716.,"The optical polarimetric campaign was conducted in early September 2008, between MJD 54710-54716."
+ Observations were made with the L6 m Perkin-Elmer telescope at Pico dos Dias Observatory of the National Astrophysics Laboratory (OPD/LN.A. Maz). using the imagine polarimeter LAGPOL in linear polarisation. moce.," Observations were made with the 1.6 m Perkin-Elmer telescope at Pico dos Dias Observatory of the National Astrophysics Laboratory (OPD/LNA, Brazil), using the imaging polarimeter IAGPOL in linear polarisation mode."
+ band images in the V. It and Ll filters were taken on all nights but the last of the campaign.," Multi-band images in the V, R and I filters were taken on all nights but the last of the campaign."
+ The configuration of the polarimeter provides simultaneous measurements of the ordinarv and extra-ordinary ανν. allowing us to perform observations under non-ideal atmospheric conditions. since any atmospheric contributions will allect both rays equally: additionally. any sky contribution is expected. to cancel out in the process.," The configuration of the polarimeter provides simultaneous measurements of the ordinary and extra-ordinary rays, allowing us to perform observations under non-ideal atmospheric conditions, since any atmospheric contributions will affect both rays equally; additionally, any sky contribution is expected to cancel out in the process."
+ Photometric Dux measurements are obtained simultaneously with the polarimetric ones., Photometric flux measurements are obtained simultaneously with the polarimetric ones.
+ Standard. polarisation stars from. Smithctal.(1991) and hector&Perlman(2008) were usec for calibration., Standard polarisation stars from \cite{b22b} and \cite{b22rp} were used for calibration.
+ Sinele polarisation images were integrated. [rom S 45 s exposures. each at a different position of the polarimetric wheel.," Single polarisation images were integrated from 8 $\times$ 45 s exposures, each at a different position of the polarimetric wheel."
+ A precision of better than in the polarisation degree was achieved., A precision of better than in the polarisation degree was achieved.
+ Phe temporal resolution of consecutive measurements in the It band was of the order of S-10 min. whereas V and L images were taken at the beginning and end. of each. night to monitor the spectral evolution of the source.," The temporal resolution of consecutive measurements in the R band was of the order of 8-10 min, whereas V and I images were taken at the beginning and end of each night to monitor the spectral evolution of the source."
+ Data reduction was mace with a specially developed analvsis package for LNA polarimetric data. PCCDPACI (Perevra2000).," Data reduction was made with a specially developed analysis package for LNA polarimetric data, PCCDPACK \citep{b19}."
+. Figure 1 shows the 1t band light-curve for the total flux. polarisation fraction ancl electric vector. position angle (EVRA) for all six nights of the optical campaign.," Figure 1 shows the R band light-curve for the total flux, polarisation fraction and electric vector position angle (EVPA) for all six nights of the optical campaign."
+ The data presented in this figure (see also Table 2 at the end of the paper) represents the directly. observed: quantities. not corrected for the unpolarised contribution of the stellar continuum.," The data presented in this figure (see also Table 2 at the end of the paper) represents the directly observed quantities, not corrected for the unpolarised contribution of the stellar continuum."
+ For the remainder of the analysis. Hux estimates as quoted in Dominiciοἱal.(2006) were used to subtract the host galaxy contribution to the total emission.," For the remainder of the analysis, flux estimates as quoted in \cite{dominici} were used to subtract the host galaxy contribution to the total emission."
+ The source was. observed. for three (ο six. hours during each night with a minimum temporal resolution in the It band of ~ 10 min. resulting in a week of well sampled intranight light-curves.," The source was observed for three to six hours during each night with a minimum temporal resolution in the R band of $\sim$ 10 min, resulting in a week of well sampled intranight light-curves."
+ The overall (ux behaviour is qualitatively distinct from the changes in the polarisation properties of the emission. as noted before by. C'ourvolsieretal.(1995). and. Tommasietal.(2001) or this same object.," The overall flux behaviour is qualitatively distinct from the changes in the polarisation properties of the emission, as noted before by \cite{cour} and \cite{tommasi} for this same object."
+ Flux variability is dominated by intranight activity. superimposed. on a baseline level which increases towards the end of the campaign and is in agreement with the measurcments from the ATOAL telescope presented. in Aharonianetal.(2009).," Flux variability is dominated by intranight activity, superimposed on a baseline level which increases towards the end of the campaign and is in agreement with the measurements from the ATOM telescope presented in \cite{b1c}."
+. A Lomb-Scargle power spectrum analvsis (Scarele1982) reveals that the Lux microvariabilitv is cleseribable as random Buctuations. with minimum variability timescales < 1 hr. limited. by the sampling of the lehteurve.," A Lomb-Scargle power spectrum analysis \citep{b22} reveals that the flux microvariability is describable as random fluctuations, with minimum variability timescales $<$ 1 hr, limited by the sampling of the lightcurve."
+ Although presenting some intranight activity. the temporal behaviour of the polarised Εαν was dominated by inter-night variations with larger relative amplitude. than those of the unpolarisecl Εαν. varving by a factor of 3 during the campaign.," Although presenting some intranight activity, the temporal behaviour of the polarised flux was dominated by inter-night variations with larger relative amplitude than those of the unpolarised flux, varying by a factor of 3 during the campaign."
+ The host-corrected polarisation degree varied smoothly. between along the six nights of observations. within the range typically registered. for the source and similar to those seen for the radio core.," The host-corrected polarisation degree varied smoothly between along the six nights of observations, within the range typically registered for the source and similar to those seen for the radio core."
+ A very similar “oscillatory” behaviour for the polarisation fraction can be seen in the optical lighteurves of Courvoisierοἱal. (1995).. but the behaviour of the polarisation vector is very distinct at both epochs.," A very similar “oscillatory” behaviour for the polarisation fraction can be seen in the optical lightcurves of \cite{cour}, but the behaviour of the polarisation vector is very distinct at both epochs."
+ Visual inspection of the light-curves shows that. the otal photometric variability cannot be explained by variations in the polarisecl [lux alone., Visual inspection of the light-curves shows that the total photometric variability cannot be explained by variations in the polarised flux alone.
+ Subtraction of the rolarised flux. from the photometric light-curves leave residual variability both in the intranight and the long-term lux variations., Subtraction of the polarised flux from the photometric light-curves leave residual variability both in the intranight and the long-term flux variations.
+ Conversely. dilution of a constant polarised component on a variable. unpolarised background. cannot account for the observed. behaviour of the polarisation degree. which changes in an uncorrelatecl fashion with respect to the total Dux.," Conversely, dilution of a constant polarised component on a variable, unpolarised background cannot account for the observed behaviour of the polarisation degree, which changes in an uncorrelated fashion with respect to the total flux."
+ Throughout our observations the EVPA underwent a quasi-lincar counter-clockwise rotation of about 407. at a rate of & 7 per day.," Throughout our observations the EVPA underwent a quasi-linear counter-clockwise rotation of about $40^\circ$, at a rate of $\approx$ $^\circ$ per day."
+ The lack of correlation between the smooth. long-term evolution of the polarisation parameters and the [ux behaviour is a common property of BL Lacs (Qianetal.1991). in optical and must be explained.," The lack of correlation between the smooth, long-term evolution of the polarisation parameters and the flux behaviour is a common property of BL Lacs \citep{qian} in optical and must be explained."
+around (he extrapolation- line. seen in. panel (a)n isH not visible∙∙ in∙ (his∙ plot any more.,around the extrapolation line seen in panel (a) is not visible in this plot any more.
+" ↴The 47> curves are given in Figure 2(a) for 13, and for d¥/d(O/IL) (b) which is discussed below.", The $\chi^2$ curves are given in Figure 2(a) for $Y_p$ and for $dY/d({\rm O/H})$ (b) which is discussed below.
+ The absorption widths derived from our fit have generally large errors. 250%.," The absorption widths derived from our fit have generally large errors, $\approx 50$."
+. The mean absorption equivalent widths are eji=0.40220.31., The mean absorption equivalent widths are $a_{\rm HeI}=0.40\pm0.31$.
+ There are 6 II] reeions for which the central values of μοι are negative., There are 6 HII regions for which the central values of $a_{\rm HeI}$ are negative.
+ The four among six are the cases for which the [its without stellar absorption gave unusually good 47., The four among six are the cases for which the fits without stellar absorption gave unusually good $\chi^2$.
+ The only cases that are mareginally acceptable are with HS10284-3843 (47=3.6) and. Mrk35 (4?=6.1). (, The only cases that are marginally acceptable are with HS1028+3843 $\chi^2=3.6$ ) and Mrk35 $\chi^2=6.1$ ). (
+In the analvsis with (joy=O. LO svstems have 4? greater than 6.1.),"In the analysis with $a_{\rm HeI}=0$, 10 systems have $\chi^2$ greater than 6.1.)"
+ OS04 presented their analvsis for seven III regions common in our sample., OS04 presented their analysis for seven HII regions common in our sample.
+ A comparison V.1i0ws Chat our Y is always consistent with theirs within one standard deviation. although je errors are significantlv larger in OSO4 due to larger uncertainties in (he temperature of je plasma.," A comparison shows that our Y is always consistent with theirs within one standard deviation, although the errors are significantly larger in OS04 due to larger uncertainties in the temperature of the plasma."
+ We find other parameters. such as temperature. are also consistent.," We find other parameters, such as temperature, are also consistent."
+ Figure 3 shows the average absorption equivalent width that varies with the heavy Tοement abundance., Figure 3 shows the average absorption equivalent width that varies with the heavy element abundance.
+ Although individual data show a rather large scatter. the trend is clear in ciis binned plot: the stellar absorption effect is more important [or metal poor LLregions!.," Although individual data show a rather large scatter, the trend is clear in this binned plot: the stellar absorption effect is more important for metal poor HII."
+. ince the oxvgen abundance shows a tight anticorrelation with temperature of the plasma. is figure is also interpreted as showing the correlation of the stellar absorption with the plasma temperature: Hieher the temperature. more (he important. absorption.," Since the oxygen abundance shows a tight anticorrelation with temperature of the plasma, this figure is also interpreted as showing the correlation of the stellar absorption with the plasma temperature: Higher the temperature, more the important absorption."
+ This is the systematic (rend that largely moclilies the extrapolation of the helium abundance to the zero metallicity., This is the systematic trend that largely modifies the extrapolation of the helium abundance to the zero metallicity.
+ In this connection another interesting quantity is dY/d(O/II). ie.. the increment of helium per heavy element production.," In this connection another interesting quantity is $dY/d({\rm O/H)}$, i.e., the increment of helium per heavy element production."
+ The [T04 results give dYα(Ο1Η)=82415 ( or dY/dZ=4.540.8. which is consistent with the final result [TO4 quoted. 1.2) from a sample of seven ILI] regions.," The IT04 results give $dY/d({\rm O/H})=82\pm 15$ ( or $dY/dZ=4.5\pm0.8$, which is consistent with the final result IT04 quoted, $dY/dZ=3.7 \pm 1.2$ ) from a sample of seven HII regions."
+ Our analvsis with αμα set equal to zero gives dY/d(O/II)=86418 or dV/dZ—4.71.0., Our analysis with $a_{\rm HeI}$ set equal to zero gives $dY/d({\rm O/H})=86 \pm 18$ or $dY/dZ=4.7 \pm 1.0$.
+ On the other hand. with stellar absorption we obtain dY/d(O/II)=20425 or dY/dZ=1.1-- 1.4. which is a clastic decrease.," On the other hand, with stellar absorption we obtain $dY/d({\rm O/H})=20 \pm 25$ or $dY/dZ=1.1 \pm 1.4$ , which is a drastic decrease."
+ We remark that [TO4 use only 3 lines in deriving the helium abundance. dropping A389. À4026 and AT065. with the anticipation that the other 3 bright lines are less affected by stellar absorption. while (hev use A3889 and AT065 (o constrain other parameters.," We remark that IT04 use only 3 lines in deriving the helium abundance, dropping $\lambda 3889$ , $\lambda 4026$ and $\lambda7065$, with the anticipation that the other 3 bright lines are less affected by stellar absorption, while they use $\lambda 3889$ and $\lambda7065$ to constrain other parameters."
+ We also (ried to drop the three lines when we caleulate Y. but the results remain unchangedfrom the full6 line analysis.," We also tried to drop the three lines when we calculate $Y$ , but the results remain unchangedfrom the full6 line analysis."
+the columu density of He at every position. not ΟΙ) and w(H)x7 separately because these two parameters. like another pair of parameters. n(H») and b. are degenerate iu the parameter space.,"the column density of $_{2}$ at every position, not $_0$ and $n$ $\times$$\tau$ separately because these two parameters, like another pair of parameters, $n$ $_{2}$ ) and $b$, are degenerate in the parameter space."
+ Iucreasing either of them will raise the resultant OPR of the gas., Increasing either of them will raise the resultant OPR of the gas.
+ Thus the confideuce intervals are wide for these two parameters. especially for n(H)xz.," Thus the confidence intervals are wide for these two parameters, especially for $n$ $\times$$\tau$."
+ The much larger uucertaiuties iu the line iutensities at a single pixel. compared. with the errors in the map-averaged inteusities. make the derived OPBRo aud 5((H)x7 even more poorly coustrained and uureliable.," The much larger uncertainties in the line intensities at a single pixel, compared with the errors in the map-averaged intensities, make the derived $_0$ and $n$ $\times$$\tau$ even more poorly constrained and unreliable."
+ That is why we show the averaged OPR maps instead., That is why we show the averaged OPR maps instead.
+ From Table 1 aud Figures 7 — 12 we see that the Ho pure rotational emissious detected by IRS of all six sources are consistent with excitation conditious in gas with (He ~2- κα and eniperature iudex 6ου 2.3 — 3.1., From Table 1 and Figures 7 – 12 we see that the $_2$ pure rotational emissions detected by IRS of all six sources are consistent with excitation conditions in gas with $n$ $_2$ ) $\sim$ 2 – 4 $\times 10^{3}$ $^{-3}$ and temperature index $b\sim$ 2.3 – 3.1.
+ In the case of W28. WIL and 339l. the intensities of H» 5(9). lowever. suggest a deuser region with v(He) 2 — s(9)2.5 times larger. auL ineudiug the IRAC baud 2 (1.5 yan) brightness — contributed mainly by H» to S(12) 2 as well as the CO higlily-excited 'otatioual line fluxes. which are considered only in the case of ICLL3C ald HH5»£ vields an even ligher best-fit density with (Ho) ~ ο x10! .," In the case of W28, W44 and 3C391, the intensities of $_2$ S(9), however, suggest a denser region with $n$ $_2$ ) 2 – 2.5 times larger, and including the IRAC band 2 (4.5 $\mu$ m) brightness – contributed mainly by $_2$ S(9) to S(12) – as well as the CO highly-excited rotational line fluxes, which are considered only in the case of IC443C and HH54, yields an even higher best-fit density with $n$ $_2$ ) $\sim$ 1 – 4 $\times10^{4}$ $^{-3}$."
+ The latter density range is closer to estimates rol previous studies of emissious [ron various species., The latter density range is closer to estimates from previous studies of emissions from various species.
+ Suell et ((20()5) found that a preshock deusity of 3:107 7 for either a slow J-type or C-type shock in IC113 clump C can account. [or he observed H3O. CO. OH aud H» 2 jun line intensities.," Snell et (2005) found that a preshock density of $\times10^{4}$ $^{-3}$ for either a slow J-type or C-type shock in IC443 clump C can account for the observed $_2$ O, CO, OH and $_2$ 2 $\mu$ m line intensities."
+ For the other three SNRs — Wes. WII aud 3C391 — the IRS regions coincide with the locatious of the brightest 1720 MHz OH inasers. he presence of which implies the existence of clumps of OH gas at iioderate temperature 50 — 125 Is aud deusities n(Ha) ~10? oE in C shocks (Lockett et 11999: Wardle YuselbZadel 2002).," For the other three SNRs – W28, W44 and 3C391 – the IRS regions coincide with the locations of the brightest 1720 MHz OH masers, the presence of which implies the existence of clumps of OH gas at moderate temperature 50 – 125 K and densities $n$ $_2$ ) $\sim10^{5}$ $^{-3}$ in C shocks (Lockett et 1999; Wardle Yusef-Zadeh 2002)."
+ For HH»L the multi-species analysis done by Ciüannuiui et (OX JOG)iudicated that an Ls km | J-type shock with a continuous precursor aid a density v(He) ~10!1 7 matches the Hoe vibrational aud pure rotational lines. as well as the CO aud H3O emissions observed with50.," For HH54, the multi-species analysis done by Giannini et (2006) indicated that an 18 km $^{-1}$ J-type shock with a continuous precursor and a density $n$ $_2$ ) $\sim 10^{4}$ $^{-3}$ matches the $_2$ vibrational and pure rotational lines, as well as the CO and $_2$ O emissions observed with."
+ Molinari et ((2000) studied the HH 7—11 outflow emissions using ZSO aud iuterpreted tle H5. CO and Ποσο line emissions as emereineeOm from a mixture of J- aud C-type sho‘ks propagating in eas of density n(Ha) ~10! cmnoE," Molinari et (2000) studied the HH 7–11 outflow emissions using ${\it ISO}$ and interpreted the $_2$, CO and $_2$ O line emissions as emerging from a mixture of J- and C-type shocks propagating in gas of density $n$ $_2$ ) $\sim 10^{4}$ $^{-3}$."
+ The inconsistency among the best-fit deusities estimated from different molecular species. obtained in the caleulatious described above. cau be explained by te deity [fluctuations within the observed regions.," The inconsistency among the best-fit densities estimated from different molecular species, obtained in the calculations described above, can be explained by the density fluctuations within the observed regions."
+ The clouds may be composed of both moclerate cleusity gas with (Ho) 10  and deuse cores with (H3) ~10?re—109 7., The clouds may be composed of both moderate density gas with $n$ $_2$ ) $\sim 10^{3}$ $^{-3}$ and dense cores with $n$ $_2$ ) $\sim10^{5}-10^{6}$ $^{-3}$.
+ Indeed. the deusi vnmaps derived [rom the IRS Ho fluxes exhibit large variations within the small areas (~1x ) mapoed. as sliown in Figures 11— 19.," Indeed, the density maps derived from the IRS $_2$ fluxes exhibit large variations within the small areas $\sim~1^{'}\times~1^{'}$ ) mapped, as shown in Figures 14 – 19."
+ The higher critical densities for the H» 5(9) to 9(12) transitiols. die CO high-/ transitions as well as the HeO lines. which we do not utilize here. make the li neijtenslties more sensitive functions of deusity than those of the Πο URS trausitions.," The higher critical densities for the $_2$ S(9) to S(12) transitions, the CO $J$ transitions as well as the $_2$ O lines, which we do not utilize here, make the line intensities more sensitive functions of density than those of the $_2$ IRS transitions."
+ Thus. denser regions contribute more to," Thus, denser regions contribute more to"
+where je is (he screening mass specilving the Debve or screening radius rj—1/j.,where $\mu$ is the screening mass specifying the Debye or screening radius $r_D=1/\mu$.
+" Ina plasma of color-charged constituents. one expects a similar behavior. aud (his is indeed observed in lattice studies. [56]: in the QGP just above Τὸ. j( increases strongly. more than linearly. and hence ry decreases correspondingly,"," In a plasma of color-charged constituents, one expects a similar behavior, and this is indeed observed in lattice studies \cite{color-screen}: in the QGP just above $T_c$, $\mu$ increases strongly, more than linearly, and hence $r_D$ decreases correspondingly."
+ Asvmptoticallv. perturbation theory suggests pog(T)T. wilh ος} lor the strong coupling running in temperature.," Asymptotically, perturbation theory suggests $\mu \simeq g(T) T$, with $g(T)$ for the strong coupling running in temperature."
+ The range of strong interactions (hus shows a striking in-medium decrease for increasing; temperature., The range of strong interactions thus shows a striking in-medium decrease for increasing temperature.
+ Quarkonia are a special kind of hadrons. bound states of a heavy (6 or 6) quark and its antiquark.," Quarkonia are a special kind of hadrons, bound states of a heavy $c$ or $b$ ) quark and its antiquark."
+ For the evound states and the binding energies are around 0.6 and 1.2 GeV. respectively. and thus much lareer than the (vpical hadronic scale À0.2 GeV: as a consequence. they are also much smaller. with radii ο of about 0.1 and 0.2 [m.," For the ground states and the binding energies are around 0.6 and 1.2 GeV, respectively, and thus much larger than the typical hadronic scale $\Lambda \sim 0.2$ GeV; as a consequence, they are also much smaller, with radii $r_Q$ of about 0.1 and 0.2 fm."
+ The fate of such states in a quark-eluon plasma therefore depends on the relative size of the color screening radius: if rp>ro. the mediun does not really affect the heavy quark binding.," The fate of such states in a quark-gluon plasma therefore depends on the relative size of the color screening radius: if $r_D \gg r_Q$, the medium does not really affect the heavy quark binding."
+ Once rp«ro. however. the (wo heavy quarks cannot “see” each other any more and hence the bound state will melt [57]..," Once $r_D \ll r_Q$, however, the two heavy quarks cannot “see” each other any more and hence the bound state will melt \cite{MS}."
+ It is therefore expected that quarkonia will survive in a quark-elion plasma through some range of temperatures above 7... and (hen melt once T becomes large enough.," It is therefore expected that quarkonia will survive in a quark-gluon plasma through some range of temperatures above $T_c$, and then melt once $T$ becomes large enough."
+ Such behavior is in [act confirmed by finite temperature lattice QCD studies of inamedium quarkonium behavior [58].., Such behavior is in fact confirmed by finite temperature lattice QCD studies of in-medium quarkonium behavior \cite{Psi-Lat}.
+ The higher excited ceuarkonium states are less tightly bound and hence Iarger. alühough (heir binding energies are in general still larger. their radii still smaller. (han those of the usual light quark hacdrons.," The higher excited quarkonium states are less tightly bound and hence larger, although their binding energies are in general still larger, their radii still smaller, than those of the usual light quark hadrons."
+ Take the charmonium spectrum as example: the radius of the J/i((19) is about 0.2 fim. that of the \.((1P) about 0.3 fm. and that of the (5) 0.1 £m.," Take the charmonium spectrum as example: the radius of the (1S) is about 0.2 fm, that of the (1P) about 0.3 fm, and that of the (2S) 0.4 fm."
+" Since melting sels in when (he screening radius reaches (he binding radius. We expect (hat the different charmonium states have dillerent ""melting temperatures! in a quark-gluon plasma."," Since melting sets in when the screening radius reaches the binding radius, We expect that the different charmonium states have different “melting temperatures” in a quark-gluon plasma."
+ llence the spectral analvsis of inanedium quarkonium dissociation should provide a QGP thermometer [59.60].," Hence the spectral analysis of in-medium quarkonium dissociation should provide a QGP thermometer \cite{KMS,KSa}."
+ As probe. we then shoot beams of specifie charmonia/t..Né. 0)) into our medium sample and check which comes out on the other side.," As probe, we then shoot beams of specific charmonia, ) into our medium sample and check which comes out on the other side."
+ If all three survive. we have an upper limit on the temperature. and by checking al just what temperature thei. the and the are dissociated. we have a way of specifying the temperature of the medium |(60].. as illustrated in reftemp..," If all three survive, we have an upper limit on the temperature, and by checking at just what temperature the, the and the are dissociated, we have a way of specifying the temperature of the medium \cite{KSa}, as illustrated in \\ref{temp}."
+ The dissociation of quarkonium states in a deconfined medium. as compared to their survival in hadronic matter. can also be considered on a more dvnanmical level. using the as," The dissociation of quarkonium states in a deconfined medium, as compared to their survival in hadronic matter, can also be considered on a more dynamical level, using the as"
+in 7..,in \citet{dabringhausen2009a}.
+ For a4=ay 2.3. Equation (1)) is the canonical IME (??)..," For $\alpha_3 = \alpha_2 = 2.3$ , Equation \ref{eq:IMF}) ) is the canonical IMF \citep{kroupa2001a,kroupa2002a}."
+ For as<2.3. the IME is top-heavy. implying more iutermiediate-auass stars and in particular more higlh-1uass stars.," For $\alpha_3 < 2.3$, the IMF is top-heavy, implying more intermediate-mass stars and in particular more high-mass stars."
+" Tn the mass range of UCDs. riya, is nof set bv the mass of the stellar svstem. but bv the observed mass luüt for stars. maga."," In the mass range of UCDs, $m_{\rm max}$ is not set by the mass of the stellar system, but by the observed mass limit for stars, $m_{\rm max *}$."
+ Thus. ηγιοςMayaxe tor all UCDs.," Thus, $m_{\rm max}=m_{\rm max *}$ for all UCDs."
+ The actual value of Hhuax is. however. rather uncertain: Estimates range from the canonical value ομωςzz150M. (07?)— to nugas2BOQADL C. but see. ?)).," The actual value of $m_{\rm max *}$ is, however, rather uncertain: Estimates range from the canonical value $m_{\rm max *} \approx 150 \ {\rm M}_{\odot}$ \citep{weidner2004a,oey2005a} to $m_{\rm max *} \approx 300 \ {\rm M}_{\odot}$ \citealt{crowther2010a}, , but see \citealt{banerjee2011a}) )."
+ Iu this paper. Hiya:=150Mi is assumed. but note that assuniüng miuyax—300M. lustead would have little effect on the results reported here (see Section 3.3. and Figure 7)).," In this paper, $m_{\rm max *} = 150 \ {\rm M}_{\odot}$ is assumed, but note that assuming $m_{\rm max *} = 300 \ {\rm M}_{\odot}$ instead would have little effect on the results reported here (see Section \ref{sec:results} and Figure \ref{fig:comp}) )."
+ Iu the case of GCs and UCDs with LMXDs (see Section 3)). he observed huuimositv. L. is1 known to originate from stars with masses m)X LAL...," In the case of GCs and UCDs with LMXBs (see Section \ref{sec:LMXB}) ), the observed luminosity, $L$, is known to originate from stars with masses $m\lesssim 1 \ {\rm M}_{\odot}$ ."
+ This is because their stellar populations are old (77)— iux the more massive stars have already: completed their evolution.," This is because their stellar populations are old \citep{evstigneeva2007a,chilingarian2008a} and the more massive stars have already completed their evolution."
+ Being fixed bx observations. £ should however not be chauged when the IME is varied.," Being fixed by observations, $L$ should however not be changed when the IMF is varied."
+ For the IME given by Equation (1)). this can be achieved bx fiudiug & from the condition where £o is the canonical IAIF. ie. a3=2.3.," For the IMF given by Equation \ref{eq:IMF}) ), this can be achieved by finding $k$ from the condition where $\xi_{\rm can}$ is the canonical IMF, i.e. $\alpha_3=2.3$."
+ With this normalization. the uwunber deusitv of stars with i)«1AL. is the same for all values of a3. suce the normalization is set by the canonical IMP aud is therefore not affected by variations of ane," With this normalization, the number density of stars with $m<1 \ {\rm M}_{\odot}$ is the same for all values of $\alpha_3$ , since the normalization is set by the canonical IMF and is therefore not affected by variations of $\alpha_3$."
+ Tu the case of the ονταῖς of Arp 220 (sce Section [)). the light used to cstimate the star formation rate (SFR). ie. the mass of the material converted mto stars per time-unit. originates fron stars over the whole rauge of stellar masses.," In the case of the SN-rate of Arp 220 (see Section \ref{sec:SNrate}) ), the light used to estimate the star formation rate (SFR), i.e. the mass of the material converted into stars per time-unit, originates from stars over the whole range of stellar masses."
+ With theSER thereby eiven. we then normalize the IME such that the SFR remains constant when the IME is varied.," With theSFR thereby given, we then normalize the IMF such that the SFR remains constant when the IMF is varied."
+ For the IMF eiven by Equation (1)). this can be achieved by finding & from the condition With this normalization. the umuber density of stars with a<1AL. decreases with decreasing values of a3. ie. with increasing top-leaviness of the IME.," For the IMF given by Equation \ref{eq:IMF}) ), this can be achieved by finding $k$ from the condition With this normalization, the number density of stars with $m<1 \ {\rm M}_{\odot}$ decreases with decreasing values of $\alpha_3$, i.e. with increasing top-heaviness of the IMF."
+ Stellar evolution aud dvuauical evolution turu the IMF of a star cluster iuto a (finic-depenudoeut) lass function of stars and stellar remnants: the star and stellar remunant mass function. SRAIF.," Stellar evolution and dynamical evolution turn the IMF of a star cluster into a (time-dependent) mass function of stars and stellar remnants; the star and stellar remnant mass function, SRMF."
+ For a siugle-age stellay population. the connection between the IAIF aud the SRME cau be quantified by an initial-to-final mass relation for stars. nna. which can be written as where mn is the mass at which stars evolve away frou the main sequence at a eiven age (?)..," For a single-age stellar population, the connection between the IMF and the SRMF can be quantified by an initial-to-final mass relation for stars, $m_{\rm rem}$, which can be written as where $m_{\rm to}$ is the mass at which stars evolve away from the main sequence at a given age \citep{dabringhausen2009a}. ."
+ UCDs typically have ages of = 10 Cir (2?).. which iuples ms;~dAL. for them.," UCDs typically have ages of $\approx$ 10 Gyr \citep{evstigneeva2007a,chilingarian2008a}, which implies $m_{\rm to} \approx 1 \ \rm{M}_{\odot}$ for them."
+ In the present paper. Equation (1)) is used to caleulate how the mass of a modeled UCD depends on the variatio- of its IMIF (see Section 3.2.3)).," In the present paper, Equation \ref{eqMrem1}) ) is used to calculate how the mass of a modeled UCD depends on the variation of its IMF (see Section \ref{sec:variableIMF}) )."
+ Note that Equation (1)) reflects the evolution of single stars., Note that Equation \ref{eqMrem1}) ) reflects the evolution of single stars.
+ Iu a binary svsteui. the imitial mas:μα of a star that evolves into a black hole is expectcca to be ligher. so that stars with masses up to 10 AL. inav become NSs iustead of DIIs (cf. 73).," In a binary system, the initial mass of a star that evolves into a black hole is expected to be higher, so that stars with masses up to 40 ${\rm M}_{\odot}$ may become NSs instead of BHs (cf. \citealt{brown2001a}) )."
+ It is however of minor miportauce in this paper whethner a massive reningut isa NS or a BIT., It is however of minor importance in this paper whether a massive remnant is a NS or a BH.
+ Both studs of objects can become bright X-ray sources weaccreting matter from a companion star and DIIs in such binary svsteis are actually detected x excludiug that they are NSs due to their mass (?).., Both kinds of objects can become bright X-ray sources by accreting matter from a companion star and BHs in such binary systems are actually detected by excluding that they are NSs due to their mass \citep{casares2007a}.
+ Also the total mass of a GC or UCD is iot strongly affected by the mass-luüt between NSs aud BOs., Also the total mass of a GC or UCD is not strongly affected by the mass-limit between NSs and BHs.
+" Using Equation AIF) with mry=LAL. and Equation (1)) with iig,=1 M... he total mass of NSs and Ες is L2 per ceut of he total mass of the stellar svstem for a3=2.3(canonical IME) and 79.9 yer cent for ay3= 1."," Using Equation ) with $m_{\rm tr} = 1 \ {\rm M}_{\odot}$ and Equation \ref{eqMrem1}) ) with $m_{to}=1 \ {\rm M}_{\odot}$ , the total mass of NSs and BHs is 4.2 per cent of the total mass of the stellar system for $\alpha_3=2.3$(canonical IMF) and 79.9 per cent for $\alpha_3=1$ ."
+ These uunbers are altered to 3.8 per cent of the otal mass ofthe stellar svsteifor a3=2.5 and 7H.0 percent for à;=Lif the transition from NSs o DIIs is shifted from 25M... to LO AL..., These numbers are altered to 3.8 per cent of the total mass ofthe stellar systemfor $\alpha_3=2.3$ and 75.0 percent for $\alpha_3=1$if the transition from NSs to BHs is shifted from $25 \ {\rm M}_{\odot}$ to $40 \ {\rm M}_{\odot}$ .
+correlation (ic. probability that it could: arise N chance).,correlation (i.e. probability that it could arise by chance).
+" ""nThis is significan enough to be useful as a distanCO indicator: he scatter in Aly for the 21 galaxies. observed is 0.326 mag. which reduces to 0.280 mag when the values are corrected for twe CORY olect."," This is significant enough to be useful as a distance indicator: the scatter in $_R$ for the 21 galaxies observed is 0.326 mag, which reduces to 0.280 mag when the $_R$ values are corrected for the $_{EW}$ effect."
+ The slope of t regression line of Aly on COR is somewhat smaller the that for the trend in My {total kW-banc absolute magnituc vs CO for Coma cluser ellipticals (Mobasher&Jam 19990).. at 1.1350.5 mag/nm ο.," The slope of the regression line of $_R$ on $_{EW}$ is somewhat smaller than that for the trend in $_K$ (total K-band absolute magnitude) vs $_{EW}$ for Coma cluster ellipticals \cite{mo:99}, at $\pm$ 0.5 mag/nm c.f."
+ -1.6x0 mag/nm for t Coma galaxies., $\pm$ 0.7 mag/nm for the Coma galaxies.
+ However. this cillerence is not statistically. significant (70.67).," However, this difference is not statistically significant $\sim$ $\sigma$ )."
+ It is not possible to determine whether the absolute CO relations are consistent [or the various samples because of the lack of homogeneous photometry. and the consequent need for large and uncertain colour and aperture corrections.," It is not possible to determine whether the absolute $_{EW}$ relations are consistent for the various samples because of the lack of homogeneous photometry, and the consequent need for large and uncertain colour and aperture corrections."
+" Similarly. there is a strong correlation between Org and residuals (NL, ) about the relaion of Alp with strucare parameter a (LloesselSQ) (lig."," Similarly, there is a strong correlation between $_{EW}$ and residuals $_{\alpha}$ ) about the relation of $_R$ with structure parameter $\alpha$ \cite{ho:80} (Fig."
+ 5). in the sense hat ealaxies with high CO tend o be bright relative to the mean relation (correlation coellicicnt 0.60. significance )).," 5), in the sense that galaxies with high $_{EW}$ tend to be bright relative to the mean relation (correlation coefficient 0.60, significance )."
+" The residuals (dM,) are reduced from 0.243 mag to 0.195 mae by correcting for the rend with COrgi shown in Fig.", The residuals $_{\alpha}$ ) are reduced from 0.243 mag to 0.195 mag by correcting for the trend with $_{EW}$ shown in Fig.
+ 5., 5.
+ There is no ¢correlation between CO and the structure parameter à 1sell., There is no correlation between $_{EW}$ and the structure parameter $\alpha$ itself.
+ Given the trends found in Figs., Given the trends found in Figs.
+ 4 5. it is instructive to explore if these elfects could cause the putative streaming low signal detected by Laucr Postman (1994)— using the full sample of BCGs.," 4 5, it is instructive to explore if these effects could cause the putative streaming flow signal detected by Lauer Postman \shortcite{la:94} using the full sample of BCGs."
+ However. we find no significant," However, we find no significant"
+where N is the number of SNe in the sample.,where $N$ is the number of SNe in the sample.
+ A summary of these fit quality statistics for all data subsets is provided in Table 2.., A summary of these fit quality statistics for all data subsets is provided in Table \ref{tab:fitstats}.
+ We can see from Figures 3 and 4 and from Table 2 that the photometric and spectroscopic cuts are doing a good job of removing serious outliers: the y? and RMS statistics improve each time a set of suspect SNe is removed.," We can see from Figures \ref{fig:sdssHubble} and \ref{fig:snlsHubble}
+ and from Table \ref{tab:fitstats} that the photometric and spectroscopic cuts are doing a good job of removing serious outliers: the $\chi^2$ and $RMS$ statistics improve each time a set of suspect SNe is removed."
+" When spectroscopic redshift priors are included, the RMS scatter from SOFT approaches the precision that can be achieved with other light curve fitters."," When spectroscopic redshift priors are included, the RMS scatter from SOFT approaches the precision that can be achieved with other light curve fitters."
+" As a recent example, ? applied MLCS2k2 and SALT-II to approximately the same data."," As a recent example, \citet{Kessler:2009} applied MLCS2k2 and SALT-II to approximately the same data."
+" Using 103 objects from the SDSS sample (comparable to our SDSS-C subset), Kessler et ffound RM5,,=0.15 mags when using MLCS2k2."," Using 103 objects from the SDSS sample (comparable to our SDSS-C subset), Kessler et found $RMS_\mu=0.15$ mags when using MLCS2k2."
+" With 62 SNe from the SNLS sample (similar to our SNLS-B), they get RMS,,=0.24 As is to be expected, when SOFT has no spectroscopic information to use for weeding out misclassifications and for defining redshift priors, the resulting fit statistics are degraded."," With 62 SNe from the SNLS sample (similar to our SNLS-B), they get $RMS_\mu=0.24$ As is to be expected, when SOFT has no spectroscopic information to use for weeding out misclassifications and for defining redshift priors, the resulting fit statistics are degraded."
+" However, even with no spectroscopic help, SOFT is still able to produce consistent and reliable redshift and distance estimates across the entire range of z. The spectroscopy-free Hubble diagram of Figure 5 shows that the increased RMS scatter is not being driven by large systematic shifts."," However, even with no spectroscopic help, SOFT is still able to produce consistent and reliable redshift and distance estimates across the entire range of z. The spectroscopy-free Hubble diagram of Figure \ref{fig:comboHubble} shows that the increased RMS scatter is not being driven by large systematic shifts."
+ The fact that reduced x? statistics in the lower half of Table 2 are close to unity indicates that the SOFT uncertainty estimates are doing a good job of representing the true error in the z and pparameter estimates., The fact that reduced $\chi^2$ statistics in the lower half of Table \ref{tab:fitstats} are close to unity indicates that the SOFT uncertainty estimates are doing a good job of representing the true error in the z and parameter estimates.
+ Recall that our redshift, Recall that our redshift
+bursts and active ealactic nuclei (ACGNs).,bursts and active galactic nuclei (AGNs).
+ Jets in Galactic N-vav binaries such as NTEJIS50-561 evolves nmchn more rapidly than ACN jets and therefore offer a good opportunity to study the dynamical evolution of relativistic jets on time scales inaccessible for ACNs., Jets in Galactic X-ray binaries such as XTEJ1550-564 evolves much more rapidly than AGN jets and therefore offer a good opportunity to study the dynamical evolution of relativistic jets on time scales inaccessible for AGNs.
+ Although afterelows in GRBs evolves also rapidly. their cosmoloeical distances make the direct imueasureiments of he proper-motion impossible aud their dynamics can be studies only inclirectly.," Although afterglows in GRBs evolves also rapidly, their cosmological distances make the direct measurements of the proper-motion impossible and their dynamics can be studies only indirectly."
+ The discovery of the extended radio anc X-ray ciission roni the microquasary NTE J550-561 (Corber et al., The discovery of the extended radio and X-ray emission from the microquasar XTE J550-564 (Corber et al.
+ 2002: Tousick et al., 2002; Tomsick et al.
+ 2003: I&aaret et al., 2003; Kaaret et al.
+ 2003) represcuts the first detection of laree-scale relativistic jets from a Calactic dack hole candidate in both radio aud X-ravs., 2003) represents the first detection of large-scale relativistic jets from a Galactic black hole candidate in both radio and X-rays.
+ These aree-scale jets appear to arise from a relatively Όσιο ejection event and. therefore. offer a unique opportuuitv o study the large-scale evolution of au iupulsive jet.," These large-scale jets appear to arise from a relatively brief ejection event and, therefore, offer a unique opportunity to study the large-scale evolution of an impulsive jet."
+ We find that the dviiuuical evolution of the observed eastern jet from the NTE J550-561 is consistent with the well-known Sedov evolutionary phase. ting which the enerev in the jet is conserved and Rox£7.," We find that the dynamical evolution of the observed eastern jet from the XTE J550-564 is consistent with the well-known Sedov evolutionary phase, during which the energy in the jet is conserved and $R\propto t^{2/5}$."
+ The apparent superhuninal motion observed at the very carly epoch inplies that the iitial motion of the jet is at least mildly relativistic., The apparent superluminal motion observed at the very early epoch implies that the initial motion of the jet is at least mildly relativistic.
+ As more aud more ISM matter is swept up. the jet decelerates and finally transits to the non-relativistic phase.," As more and more ISM matter is swept up, the jet decelerates and finally transits to the non-relativistic phase."
+ A trans-relativistic external shock dvuamical model is shown to be able to fit the observed proper motion cata reasonably well., A trans-relativistic external shock dynamical model is shown to be able to fit the observed proper motion data reasonably well.
+ The inferred ISAT density around the value.jet iuo1510tem°. wellbelow the canonical ," The inferred ISM density around the jet is $n\sim1.5\times10^{-4}{\rm cm^{-3}}$, well below the canonical value."
+Such low deusity gains support from the inferred value 6iISMO10Sem7 around another two microquasars CRS 1915|105 and GRO J1655-I0 by IIeiuz (2002). from he fact that jets move with constant velocities (i.e. no slowing down) up to distance z0.01 pe.," Such low ISM density gains support from the inferred value $n\la 10^{-3}{\rm cm^{-3}}$ around another two microquasars GRS 1915+105 and GRO J1655-40 by Heinz (2002), from the fact that jets move with constant velocities (i.e. no slowing down) up to distance $\ga0.04$ pc."
+ As suggested x Πας (2002). this nuplies either that the sources are ocated in regions occupied by the hot ISAL plase or hat previous frequent activities of the jets have created evacuated bubbles around the sources.," As suggested by Heinz (2002), this implies either that the sources are located in regions occupied by the hot ISM phase or that previous frequent activities of the jets have created evacuated bubbles around the sources."
+ We first trv to ft the X-ray light curve of the caster jet with the cmiussion from the shocked ISM., We first try to fit the X-ray light curve of the eastern jet with the emission from the shocked ISM.
+ However. it is found that this predicts a decay too slow to fit the observations.," However, it is found that this predicts a decay too slow to fit the observations."
+" The model predicts FE,x¢(or210/10t+? diving the non-relativistic phase. while the power law fit of the N-ray flux data eives FyxtOT!"" (Iuuct et al."," The model predicts $F_\nu\propto t^{-(15p-21)/10}\sim t^{-1.2}$ during the non-relativistic phase, while the power law fit of the X-ray flux data gives $F_\nu\propto t^{-3.7\pm0.7}$ (Kaaret et al."
+ 2003)., 2003).
+ We then turu to another Likely enüssion regionthe adiabatically expanding ejecta heated by the reverse shock. quite similar to the mechanisin suggested to be responsible for the optical flash aud radio flare fromm GRD990123 (Sari Pirau 1999).," We then turn to another likely emission region—the adiabatically expanding ejecta heated by the reverse shock, quite similar to the mechanism suggested to be responsible for the optical flash and radio flare from GRB990123 (Sari Piran 1999)."
+ Different from the shocked ISM. all electrous iu the ejecta cool by adiabatic expansion. so the πιαπα οποίος of the electrons in the ejecta decreases as," Different from the shocked ISM, all electrons in the ejecta cool by adiabatic expansion, so the maximum energy of the electrons in the ejecta decreases as."
+ Ouce the characteristic svuchrotron radiation frequency of these electrons falls close to the N-rav. baud. the N-rav. flux from the ejecta should decay quite rapidly (drops exponentially with time) since then.," Once the characteristic synchrotron radiation frequency of these electrons falls close to the X-ray band, the X-ray flux from the ejecta should decay quite rapidly (drops exponentially with time) since then."
+ Usine this model. we fitted both the N-rax light curve data aud the energy spectra on 1 June 2000 of the eastern jet successfully (see Figures Land 5).," Using this model, we fitted both the X-ray light curve data and the energy spectrum on 1 June 2000 of the eastern jet successfully (see Figures 4 and 5)."
+ The western (receding) jet from NTE J1550-56 πας detected iu radio and N-avs in 2002 while archivalChandra data on this source from June. August and September 2000 only eive upper Bits.," The western (receding) jet from XTE J1550-564 was detected in radio and X-rays in 2002 while archival data on this source from June, August and September 2000 only give upper limits."
+ The non-detection in 2000 is consistent with the deduction that the westeru source is the receding jet., The non-detection in 2000 is consistent with the deduction that the western source is the receding jet.
+ Culike the sinoothlv decaving eastern jet. the western jet brighteus at the late time.," Unlike the smoothly decaying eastern jet, the western jet brightens at the late time."
+ The brightening might be caused by the inhomogencitics in the ISAL (saaret et al., The brightening might be caused by the inhomogeneities in the ISM (Kaaret et al.
+ 2003) or internal shocks produced by a faster jet overtaking a slower one (INaiscr et al., 2003) or internal shocks produced by a faster jet overtaking a slower one (Kaiser et al.
+ 2000). and ποσα further careful study.," 2000), and need further careful study."
+" The western jet moved by 1,52-E0.13 are sec between 11 March aud 19 June 2002 with alnean apparent speed significantly less than the average apparent speed from 1998 to carly 2002.", The western jet moved by $0.52\pm0.13$ arc sec between 11 March and 19 June 2002 with a mean apparent speed significantly less than the average apparent speed from 1998 to early 2002.
+ Interestingly. we find that the decay of the N-rav. flux of the western jet οσοι March aud June 2002 is also consistent with that xedieted by the reverse shock cunission #186 foy p=2.32 of the western jet.," Interestingly, we find that the decay of the X-ray flux of the western jet between March and June 2002 is also consistent with that predicted by the reverse shock emission $F_\nu\propto t^{-4/5p}\propto t^{-1.86}$ for $p=2.32$ of the western jet."
+ Besides the relativistically movine. jets your NTE J1550-561. lavee-scale N-vaydecelerating jets and radio obes up to ~10arcinin size have been observed from $$133 (Brinkmann et al.," Besides the relativistically moving, decelerating jets from XTE J1550-564, large-scale X-ray jets and radio lobes up to $\sim 40~
+{\rm arcmin}$ size have been observed from SS433 (Brinkmann et al."
+ 1996: Dubner etvunination al., 1996; Dubner et al.
+1998)., 1998).
+ The radiation is suggested to come from the te shock which results from the interaction of the mass outilow with the nebula W50., The radiation is suggested to come from the termination shock which results from the interaction of the mass outflow with the nebula W50.
+ Recently reheating of barvonic uaterial iu X-ray romjets of SS133 is interred to take place within ~10!cm the core. based on the observed iron eniissionu lines (Mieliaui. Feuder Méuudez 2002).," Recently, reheating of baryonic material in X-ray jets of SS433 is inferred to take place within $\sim 10^{17}{\rm cm}$ from the core, based on the observed iron emission lines (Migliari, Fender Ménndez 2002)."
+ We link that exterual shock is a possible mechanisin for such reheating., We think that external shock is a possible mechanism for such reheating.
+ As sugeested for TeV uecutrino ciission from wicroquasar jets (through internal shocks) bv Levinsou Wasiman (2002). external shocks of microquasar jets nay also accelerate the protons in both the shocked ISAL aud the shocked ejecta.," As suggested for TeV neutrino emission from microquasar jets (through internal shocks) by Levinson Waxman (2002), external shocks of microquasar jets may also accelerate the protons in both the shocked ISM and the shocked ejecta."
+ So. they are also poteutial sources of cosniüc-ravs (Ileiuz Suuvaev 2002). eucreyv neutrinos and high-cucrey gauunuia-ravs.," So, they are also potential sources of cosmic-rays (Heinz Sunyaev 2002), high-energy neutrinos and high-energy gamma-rays."
+ Iu sunumuary. we developed a model for the dynamical evolution and radiation of the large-scale N-raw jets roni the nücroquasar NTE J1550-561 analogous to the external shock model for GRB afterelows.," In summary, we developed a model for the dynamical evolution and radiation of the large-scale X-ray jets from the microquasar XTE J1550-564 analogous to the external shock model for GRB afterglows."
+ In this model. he observed jet cussion is due to interaction between he jets and external ISML," In this model, the observed jet emission is due to interaction between the jets and external ISM."
+ Future continuous. follow-up imultiavaveleueth observations of new ejection events roni nicroqusass up to the significant deceleration pliase should provide more valuable insights iuto the nature aud physical condition (e$. shock aud particle acceleration uivsies) of relativistie jets.," Future continuous, follow-up multi-wavelength observations of new ejection events from microqusasrs up to the significant deceleration phase should provide more valuable insights into the nature and physical condition (e.g. shock and particle acceleration physics) of relativistic jets."
+ Owing to the proximity of the Galactic X-ray binaries. further studies ou them also offer an exciting wav for a better understanding of relativistic jets seen clsewhere in the Universe.," Owing to the proximity of the Galactic X-ray binaries, further studies on them also offer an exciting way for a better understanding of relativistic jets seen elsewhere in the Universe."
+with protons. but including electron collisions (aud asstuning that Πρ—n; aud T.—10! I). we have up=9.31x10 Loan? 1 (Seaton1955).,"with protons, but including electron collisions (and assuming that $n_e = n_i$ and $T = 10^4$ K), we have $C_{sp} = 5.31
+\times 10^{-4}$ $^3$ $^{-1}$ \citep{sea55}."
+. Since AT is much greater than the energy separations between the 2s aud 2p states. tlie reverse rate Is determined by just the ratio of statistical weights. Le. Sinilar conskleratious apply to the 2p states.," Since $kT$ is much greater than the energy separations between the $2s$ and $2p$ states, the reverse rate is determined by just the ratio of statistical weights, i.e. Similar considerations apply to the $2p$ states."
+ The fraction of recombinatious reaching 2p is (1—f). [, The fraction of recombinations reaching $2p$ is $(1-f)$. [
+Et is worth noting that processes such as Lyman-3 absorption. which ultimately produce the 25 state through 15—3p2s. have already been takeu into account in the caleulation of f. (Spitzer 1951)..],"It is worth noting that processes such as $\beta$ absorption, which ultimately produce the $2s$ state through $1s\rightarrow 3p\rightarrow 2s$, have already been taken into account in the calculation of $f$ \citep{spi51}. .]"
+" Radiative decay occurs through both spontaneous Lyiman-o emission (with Ao,=6.25x105 1). as well as stimulated emission."," Radiative decay occurs through both spontaneous $\alpha$ emission (with $A_{21} = 6.25 \times 10^8$ $^{-1}$ ), as well as stimulated emission."
+ We must also include radiative excitation from the ground state via Lymanu-a absorption., We must also include radiative excitation from the ground state via $\alpha$ absorption.
+ Includiug collisions which couple to the 2s states. the rate equation for 2p then is (Ls) ∖∖↽∐≺↵↕⋅≺↵∣∣∣∕⋜↕∐≼⇂∣∣∣∕↙⋜⋃⋅≺↵↕∐≺↵≺⇂↩∐⊳∖↓⋃≺↵⊳∖∩↥⋜↕↕∩⋯⊳∖⋯↕∐≺↵↥⋅↣⇀⋜↕∐≺⇂−≻∣↗⊳∖↕⋜⋯↵⊳∖↥↽≻≺↲↥⋅↕⋅⋜↕≺∐⋜↕↥∖⊽≺↵↥∩∢∙∐⊽∖⇁(2p) ⋅⋅ ⋅ ⋅ ↽ ⋅∣ ⋅ ⋅ ↕∐↕↩↓⋅∖⊽⋜↕↥⋅∐↕↩⋜↕⊳∖⋯⋅≺↵≺⊔∐∐⋅≺↵≺↽↓⋯↵∐∢∙⊽∖⊽⋃∐∐⊳∖∶⋜↕∐≺⇂∣∣∣∕∕↥⊳∖↕∐≺↲↥↽≻∐∩↕∩↥⊔⇂≺↵∐⊳∖∐⊽∖⇁↥↽≻≺↵↕⋅↥∎↕⋅≺↵≺↽↓⋯↲∐∢∶∖⇁∐∐≺↵↕⋅∖⇁⋜↕↥⋖∫∐↕ ∢↝⋯⊽∎∐∠↓⋝⋅," Including collisions which couple to the $2s$ states, the rate equation for $2p$ then is where $n_\nu^{(1s)}$ and $n_\nu^{(2p)}$ are the densities of atoms in the $1s$ and $2p$ states per radial velocity interval, measured in frequency units; and $n_{\nu^\prime}$ is the photon density per frequency interval (in $^{-3}$ $^{-1}$ )."
+↓∖⊽∩∖∖↽⋅ where vp is the Lyman-a frequency. and L(r—vv’) is the Lorentz line profile.," Now, where $\nu_L$ is the $\alpha$ frequency, and $L(\nu - \nu^\prime)$ is the Lorentz line profile."
+ For a therma eas. where (Av)p is the Doppler width.," For a thermal gas, where $(\Delta\nu)_D$ is the Doppler width."
+ For T=10! I. (.Nv)p=1.29xLOM Hz.," For $T = 10^4$ K, $(\Delta\nu)_D = 1.29 \times 10^{11}$ Hz."
+ Clearly therma widths will completely dominate the natural (Lorentz) line width: of Lymau-a. aud thus we cai replace L(r—/) by the Dirac delta function.," Clearly thermal widths will completely dominate the natural (Lorentz) line width of $\alpha$, and thus we can replace $L(\nu - \nu^\prime)$ by the Dirac delta function."
+ We then obtain for the absorption term: Of course. a similar calculation eau be carried out for the stimulated emission term.," We then obtain for the absorption term: Of course, a similar calculation can be carried out for the stimulated emission term."
+ Because the nebula is optically thick in Lyiuau-a. n; must be consklered carefully.," Because the nebula is optically thick in $\alpha$, $n_\nu$ must be considered carefully."
+ Were escape from the nebula the primary removal mechanism. then a steady state would result in which pliotons," Were escape from the nebula the primary removal mechanism, then a steady state would result in which photons"
+our cases of resonant motion between the irregular satellites of Saturn: The authors computed. the orbital evolution of hese satellites with ai mocified version of Swift. NBods code. where the planets were integrated in. the icliocentrie reference Frame while the irregular satellites in xanetocentric frames.,"four cases of resonant motion between the irregular satellites of Saturn: The authors computed the orbital evolution of these satellites with a modified version of Swift N--Body code, where the planets were integrated in the heliocentric reference frame while the irregular satellites in planetocentric frames."
+ This dvnamical structure is similar o that of our Model 1., This dynamical structure is similar to that of our Model $1$.
+ We extracted from our simulations he data concerning the same objects studied in al. (2002).. Nesvornyetal.(2003) and Cuk&Burns(2004) and we analysed their behaviour.," We extracted from our simulations the data concerning the same objects studied in \cite{car02}, \cite{nes03} and \cite{cuk04} and we analysed their behaviour."
+ The motion of Liraq as rom our Models 1 and 2 integrated with LLJS (sec fig. 5..," The motion of Ijiraq as from our Models $1$ and $2$ integrated with HJS (see fig. \ref{ijiraq-peri},"
+ irst three panels from top left in counterclockwise direction) appears to evolve in a stable Ixozai regime with the Sun., first three panels from top left in counterclockwise direction) appears to evolve in a stable Kozai regime with the Sun.
+ The longitude of pericenter. librates around 90° with an amplitude of 430°., The longitude of pericenter librates around $90^{\circ}$ with an amplitude of $\pm30^{\circ}$.
+ Phe simulations based on Model 2 and Alodel 1 (SO bit precision) show a regular secular behaviour of the orbital elements while the simulation based on Moclel 1 (64 bit precision) shows periods in which the range of variation of the eccentricitv. shrinks by a few percent. in correspondence to a similar behaviour of the longitude of pericenter (see fie., The simulations based on Model $2$ and Model $1$ $80$ bit precision) show a regular secular behaviour of the orbital elements while the simulation based on Model $1$ $64$ bit precision) shows periods in which the range of variation of the eccentricity shrinks by a few percent in correspondence to a similar behaviour of the longitude of pericenter (see fig.
+ ο for details)., \ref{ijiraq-pe} for details).
+ The output. obtained from Model 2 with RADAU algorithm (fis. 7)), The output obtained from Model $2$ with RADAU algorithm (fig. \ref{ijiraq-pei}) )
+ is instead significantly dillerent., is instead significantly different.
+ [iraq is in a stable Wozai regime for the first 6«10 vears. then it experiences a change in the secular behaviour of both eccentricity and inclination and. finally. the longitude of pericenter zo. circulates for," Ijiraq is in a stable Kozai regime for the first $6 \times 10^{7}$ years, then it experiences a change in the secular behaviour of both eccentricity and inclination and, finally, the longitude of pericenter $\varpi$ circulates for"
+6€770035|4348. the prominent νοz outlier. is included in this calculation. then the dispersion becomes c0.75 mag.,"6C**0935+4348, the prominent $K-z$ outlier, is included in this calculation, then the dispersion becomes $\simeq 0.75$ mag."
+ The estimated redshift distribution. for the entire. 6€ sample is presented in Fig S.., The estimated redshift distribution for the entire 6C** sample is presented in Fig \ref{fig:histzest}.
+ Phe median estimated redshift is z&LG., The median estimated redshift is $z \simeq 1.6$.
+ Ht is informative to compare these results with the redshift distribution of the 6C* sample (Jarvis et. al., It is informative to compare these results with the redshift distribution of the 6C* sample (Jarvis et al.
+ 2001b) which. as mentioned before. has been filtered in a very similar way.," 2001b) which, as mentioned before, has been filtered in a very similar way."
+ The 6C* sample has a median redshift of zo19 and a redshift cüstribution skewed towards z22., The 6C* sample has a median redshift of $z \simeq 1.9$ and a redshift distribution skewed towards $z > 2$.
+ This is à direct result of the filtering criteria applied to it., This is a direct result of the filtering criteria applied to it.
+ Although the distribution of 6€** sources does have a tail to high. redshift. the ereat majority of sources are ab22. dn clear contrast to what should. be expected.," Although the distribution of 6C** sources does have a tail to high redshift, the great majority of sources are at, in clear contrast to what should be expected."
+ However. and in light of Section. 4.. we have to consider what is likely to be. the major source of bias in our redshift estimation: the presence of quasars in the sample.," However, and in light of Section \ref{sec:comparison}, we have to consider what is likely to be the major source of bias in our redshift estimation: the presence of quasars in the sample."
+" ""These sources will be the ones which will be skewing the distribution the most. towards lower redshifts.", These sources will be the ones which will be skewing the distribution the most towards lower redshifts.
+ “Phe redshift’ estimates of the three quasars with unresolved identifications (6C**O714|4616. ο**0922|4216 and. ος*L036|4721). for which we have a spectroscopic redshift. are under-estimated by an average [actor of 0.4.," The redshift estimates of the three quasars with unresolved identifications (6C**0714+4616, 6C**0922+4216 and 6C**1036+4721), for which we have a spectroscopic redshift, are under-estimated by an average factor of 0.4."
+ Phe major cause for concern is that there are quasars for which we do not have spectroscopic information., The major cause for concern is that there are quasars for which we do not have spectroscopic information.
+ On the basis of imaging alone. we assume that the further following sources are quasars (given that they have bright unresolved identifications: Paper D): ος**0849|4658. 6C7*1003|4827. *1052]4349. 6C**1056]5730 and GCPTTBS|3803.," On the basis of imaging alone, we assume that the further following sources are quasars (given that they have bright unresolved identifications; Paper I): 6C**0849+4658, 6C**1003+4827, 6C**1052+4349, 6C**1056+5730 and 6C**1138+3803."
+ This leaves us with ninequasars'.. eight of which have redshift estimates in the range1.," This leaves us with nine, eight of which have redshift estimates in the range."
+0.. ‘These alone represent a significant source of the bias towards Iow-recdshift which we witness in Fig δ.., These alone represent a significant source of the bias towards low-redshift which we witness in Fig \ref{fig:histzest}.
+ Removing the quasars from the distribution results in a median estimated redshift of z21.7., Removing the quasars from the distribution results in a median estimated redshift of $z \simeq 1.7$.
+ This is similar to what is found for 6C* and much higher than the median redshift of similar. unfiltered samples (e.g. the TORS sample. at the same I[lux-densitv limit. with z 1.1).," This is similar to what is found for 6C* and much higher than the median redshift of similar, unfiltered samples (e.g. the 7CRS sample, at the same flux-density limit, with $z \approx 1.1$ )."
+ This result confirms the cllicieney of the filtering. criteria. used. for 6C**. in excluding low-recshift sources.," This result confirms the efficiency of the filtering criteria, used for 6C**, in excluding low-redshift sources."
+ Furthermore. we recall that some sources mav also have their redshifts. svstematically uncer-estimated. due to emission-line contamination to their magnitudes (Sections 3.2. and 4)).," Furthermore, we recall that some sources may also have their redshifts systematically under-estimated, due to emission-line contamination to their magnitudes (Sections \ref{sec:emissionline} and \ref{sec:comparison}) )."
+ Thus. it is quite plausible that the real redshift distribution has a slightly higher median recdshift.," Thus, it is quite plausible that the real redshift distribution has a slightly higher median redshift."
+ The fraction of quasars in the 6C sample (9 quasars out of GS objects) is higher than that in the 6€57 sample., The fraction of quasars in the 6C** sample (9 quasars out of 68 objects) is higher than that in the 6C* sample.
+ There are only. two quasars out of the 29 6C% objects (Jarvis ct al., There are only two quasars out of the 29 6C* objects (Jarvis et al.
+ 2001b)., 2001b).
+ Although the clillerence is. barely significant given the small numbers involved. it is possibly inlluenced by small dillerences in. the selection criteria. such as the tighter size constraint or. more importantly. the lower frequency. spectral index cut.," Although the difference is barely significant given the small numbers involved, it is possibly influenced by small differences in the selection criteria, such as the tighter size constraint or, more importantly, the lower frequency spectral index cut."
+ For 607. the steep spectral index constrain is evaluated: between MMIIz and GGlIIz. thus excluding objects with prominent Iat-spectrum Cores. Le. quasars.," For 6C*, the steep spectral index constrain is evaluated between MHz and GHz, thus excluding objects with prominent flat-spectrum cores, i.e. quasars."
+ Because the 6€ spectral index cut goes up to only L4AGGILz. this elect ds. less pronounced.," Because the 6C** spectral index cut goes up to only GHz, this effect is less pronounced."
+corresponds to 180 kin f which is narrower (han the FWHM that would be inferred from the stellar velocity dispersion (248 km |. or a EWIIM of 584 kin 4).,"corresponds to 180 km $^{-1}$, which is narrower than the FWHM that would be inferred from the stellar velocity dispersion (248 km $^{-1}$, or a FWHM of 584 km $^{-1}$ )."
+" NGC 3923: This galaxy has significant extinction and a moderately large III column (6.3x 107"" 7). although the Galactic Hs lines are not especially strong."," NGC 3923: This galaxy has significant extinction and a moderately large HI column $\times$ $^{20}$ $^{-2}$ ), although the Galactic $_{2}$ lines are not especially strong."
+ No emission lines are seen (Figure 11)., No emission lines are seen (Figure 11).
+ The weaker OVI line is in an uncontaminated part of the spectrum. while the stronger OVI line would coincide with a weak Galactic Ls absorption line ancl is adjacent to strong airglow lines.," The weaker OVI line is in an uncontaminated part of the spectrum, while the stronger OVI line would coincide with a weak Galactic $_{2}$ absorption line and is adjacent to strong airglow lines."
+ NGC 4125: The stellar continuum is very weak in (his galaxy. despite one of the longest observations in the sample (Figure 12).," NGC 4125: The stellar continuum is very weak in this galaxy, despite one of the longest observations in the sample (Figure 12)."
+ The continuum is so poorly defined (hat no Galactic absorption lines can be identified., The continuum is so poorly defined that no Galactic absorption lines can be identified.
+ The night-onlv data of the total exposure time) sel vields stricter upper limits to the OVI emission., The night-only data of the total exposure time) set yields stricter upper limits to the OVI emission.
+ No emission features [rom any lines are detected. despite this galaxy. being listed as a LINER and classified as EG pec.," No emission features from any lines are detected, despite this galaxy being listed as a LINER and classified as E6 pec."
+ NGC 4374: This ellipcal in the Virgo cluster has radio lobes ancl (his is classified as a LINER., NGC 4374: This elliptical in the Virgo cluster has radio lobes and this is classified as a LINER.
+ Despite being at high Galactic latitude = 747). it has moderate absorption ancl the spectrum (Figure 13) shows clear evidence for Galactic Hs absorption.," Despite being at high Galactic latitude = $^{\circ}$ ), it has moderate absorption and the spectrum (Figure 13) shows clear evidence for Galactic $_{2}$ absorption."
+ The day spectrum had particularly strong airglow lines. so we used the night-onlv spectrum. which had only 3.7 ksec. vel it reveals a strong OVI AI032 al the galaxy redshift (also present ad (he same level in the dav-night spectrum): (he red side of the line may be partly absorbed bv the Galactic CII AI036 line.," The day spectrum had particularly strong airglow lines, so we used the night-only spectrum, which had only 3.7 ksec, yet it reveals a strong OVI $\lambda$ 1032 at the galaxy redshift (also present at the same level in the day+night spectrum); the red side of the line may be partly absorbed by the Galactic CII $\lambda$ 1036 line."
+ The weaker OVI line lies between (wo airglow lines and there is an instrumental feature near 1044.A. but a line half of the strength of the OVI ALI032 line is consistent with the level of the spectrum above (hat expected from the stars.," The weaker OVI line lies between two airglow lines and there is an instrumental feature near 1044, but a line half of the strength of the OVI $\lambda$ 1032 line is consistent with the level of the spectrum above that expected from the stars."
+ The line width of 130 km ! (FWIIM) is only about one-fifth the velocity dispersion of the stars. an equivalent FWILM of 674 kan |. although we cannot rule out that the line has a broader base.," The line width of 130 km $^{-1}$ (FWHM) is only about one-fifth the velocity dispersion of the stars, an equivalent FWHM of 674 km $^{-1}$, although we cannot rule out that the line has a broader base."
+" NGC 4406: Also known as M86. this lies in the central part of the cluster. 1.2° west of M87. and with a modest IHE column (2.6x 107"" 7)."," NGC 4406: Also known as M86, this lies in the central part of the cluster, $^{\circ }$ west of M87, and with a modest HI column $\times$ $^{20}$ $^{-2}$ )."
+ It has a nearly radial orbit. as iis redshift (-244 kms !) differs from the svstemic velocity of the cluster by more than the cluster velocity dispersion.," It has a nearly radial orbit, as its redshift (-244 km $^{-1}$ ) differs from the systemic velocity of the cluster by more than the cluster velocity dispersion."
+ It has an elongated X-ray distribution (hat has been interpreted as stripping of the gas within NGC 4406 bv the ambient cluster medium., It has an elongated X-ray distribution that has been interpreted as stripping of the gas within NGC 4406 by the ambient cluster medium.
+ The stellar continuum is verv [aint so that it is difficult to detect anv Galactic absorption features (Figure 14)., The stellar continuum is very faint so that it is difficult to detect any Galactic absorption features (Figure 14).
+ The stronger OVI line would be shifted to 1031.1.A. and a possible emission feature in this low S/N spectrum occurs at that location.," The stronger OVI line would be shifted to 1031.1, and a possible emission feature in this low S/N spectrum occurs at that location."
+ Although (his is the broadest and most significant enission feature. il is less than a 36 detection.," Although this is the broadest and most significant emission feature, it is less than a $\sigma$ detection."
+ The weaker OVI line is not detected. nor are anv other emission lines.," The weaker OVI line is not detected, nor are any other emission lines."
+ NGC 4472: Known as MA9. (hislies in the center of the southern erouping of the Virgo," NGC 4472: Known as M49, thislies in the center of the southern grouping of the Virgo"
+mid-IR sources.,mid-IR sources.
+ It is possible that some of the AGNs are highly obscured even at 15 jan because they are viewed through a verv large dust column., It is possible that some of the AGNs are highly obscured even at 15 $\mu$ m because they are viewed through a very large dust column.
+ A mid-IR source obscured by a nearly Compton-(hick (7.=0.7) disk with Galactic dust/gas ratio would have an equatorial extinction of e5 mag (factor of 100) at 15 pam. Even in this case. any mid-IR emission above the disk might not be obscured.," A mid-IR source obscured by a nearly Compton-thick $\tau_\mathrm{e}=0.7$ ) disk with Galactic dust/gas ratio would have an equatorial extinction of $\sim 5$ mag (factor of 100) at 15 $\mu$ m. Even in this case, any mid-IR emission above the disk might not be obscured."
+ For example. there mav be a contribution from dust in the narrow-line region (NLR). above the hole in the torus (e.g.NGC1063.GallianoDocketal. 2000).. tending to make the mid-IR? emission more isotropic.," For example, there may be a contribution from dust in the narrow-line region (NLR), above the hole in the torus \citep[e.g. NGC 1068,][]{gpa05,bnm00}, tending to make the mid-IR emission more isotropic."
+ For galaxies with redshift 2<0.22. Spitzer can measure the flux at A=30 jm (rest). which should be more isotropic and less subject to extinction than the 15 jm flux.," For galaxies with redshift $z \le 0.22$ , can measure the flux at $\lambda = 30$ $\mu$ m (rest), which should be more isotropic and less subject to extinction than the 15 $\mu$ m flux."
+" Nevertheless. 9/11 mid-IR weak galaxies in this redshift range are also weak at 30 pam. with vL,(30 pm)«8x107 erg |."," Nevertheless, 9/11 mid-IR weak galaxies in this redshift range are also weak at 30 $\mu$ m, with $\nu L_\nu(30$ $\mu\mathrm{m}) <8 \times 10^{43}$ erg $^{-1}$."
+" The two exceptions are 3C 61.1 and 3C 123. with »L,(30 gon)=L2TX0.08 and 1440.2x107 erg +. respectively."," The two exceptions are 3C 61.1 and 3C 123, with $\nu L_\nu(30$ $\mu\mathrm{m})=1.27 \pm 0.08$ and $1.4\pm 0.2\times 10^{44}$ erg $^{-1}$, respectively."
+" In comparison 3C 452. the weakest mid-IR luminous NLRG in this redshift range. has vL,(30 pim)=8.63£0.09x1077 erg +."," In comparison 3C 452, the weakest mid-IR luminous NLRG in this redshift range, has $\nu L_\nu(30$ $\mu\mathrm{m})=8.63 \pm 0.09\times 10^{43}$ erg $^{-1}$."
+ Therefore. reclassifving the NLRGs by their luminosity at 30 san would only gain us an additional 2 mid-IR luminous sources.," Therefore, reclassifying the NLRGs by their luminosity at 30 $\mu$ m would only gain us an additional 2 mid-IR luminous sources."
+ This leads us to believe that most of the weak radio galaxies truly lack à powerful accretion disk., This leads us to believe that most of the mid-IR weak radio galaxies truly lack a powerful accretion disk.
+ Relatively low accretion power suggests. bul does not prove. that some FR. 1I jets may be driven by radiativelv inellicient accretion [lows or black hole spin-energv (Degelman.Blancllorcl.&Rees1934:Meier1999).," Relatively low accretion power suggests, but does not prove, that some FR II jets may be driven by radiatively inefficient accretion flows or black hole spin-energy \citep{bbr84, m99}."
+. As we mentioned above. roughly half (9/17) of the mid-IR weak NLRCGs are are classilied as LEGs with weak optical [O ΠΠ emission.," As we mentioned above, roughly half (9/17) of the mid-IR weak NLRGs are are classified as LEGs with weak optical [O ] emission."
+ The [O 11] emission in these sources may be weak because there is no strong source of UV. photous to power the NLR., The [O ] emission in these sources may be weak because there is no strong source of UV photons to power the NLR.
+ Qualitativelv similar conclusions have been drawn by other investigators (e.g..[line&Longair1979:Chiabergeetal.2000:Grimes.Rawlings.&Willott 2004).," Qualitatively similar conclusions have been drawn by other investigators \citep[e.g.,][]{hl79, ccc00,grw04}."
+.. Alternatively. the NL. maa be partly or completely obscured in LEGs (IILaasetal...2005).," Alternatively, the NLR may be partly or completely obscured in LEGs \citep{hss5}."
+. It will be important to make a more quantitative assessment of (he optical ancl IR emission line strengths. to evaluate (he extinction and determine what UV Iuminositv is necessary lo power the NLIR in mid-In weak NERGs.," It will be important to make a more quantitative assessment of the optical and IR emission line strengths, to evaluate the extinction and determine what UV luminosity is necessary to power the NLR in mid-IR weak NLRGs."
+ One of the major motivations for the radio galaxy and quasar unification (heorv is the deficit of lobe-dominant vs. core-dominant FR IE quasars (Barthel1989)., One of the major motivations for the radio galaxy and quasar unification theory is the deficit of lobe-dominant vs. core-dominant FR II quasars \citep{b89}.
+. Relativistic beaming models predict that there should be relatively more sources where the radio jet is beanmied away andthe high frequency radio core is weak., Relativistic beaming models predict that there should be relatively more sources where the radio jet is beamed away andthe high frequency radio core is weak.
+ Core fluxes αἱ 5 GlIz are tabulated for the 3CRIX catalog by Laingetal.(2003) http://www.scerr.dyvuedus.ore..., Core fluxes at 5 GHz are tabulated for the 3CRR catalog by \cite{lrl03} .
+ We identify, We identify
+llaardt for making his updated: UV. background. mocel available to us.,Haardt for making his updated UV background model available to us.
+ JSB thanks John Regan for assistance with running simulations. and Natasha Alacddox ancl Paul Llewett for helpful discussions.," JSB thanks John Regan for assistance with running simulations, and Natasha Maddox and Paul Hewett for helpful discussions."
+ Εις research was conducted in cooperation with SCl/Intel. utilising the Altix 3700 supercomputer COSMOS at the. Department of Applied Alathematies and Theoretical Physics in Cambridge., This research was conducted in cooperation with SGI/Intel utilising the Altix 3700 supercomputer COSMOS at the Department of Applied Mathematics and Theoretical Physics in Cambridge.
+ COSMOS is a Ul-CCC. facility which is supported. by HEPCE and PPARC., COSMOS is a UK-CCC facility which is supported by HEFCE and PPARC.
+ We also acknowledge support from the European Community Rescarch and Training Network “Lhe Physics of the Intergalactic Alecium’., We also acknowledge support from the European Community Research and Training Network 'The Physics of the Intergalactic Medium'.
+ JSB. MGIIL. MV and REC thank PPARC for financial support.," JSB, MGH, MV and RFC thank PPARC for financial support."
+ We also thank the referee. Alike Shull. for a very detailed and helpful report.," We also thank the referee, Mike Shull, for a very detailed and helpful report."
+terminal velocity of theIL3 absorption component was —390 kins !. while in our spectrum ibis —530 kms !.,"terminal velocity of the$\beta$ absorption component was $-$ 390 km $^{-1}$, while in our spectrum it is $-$ 530 km $^{-1}$."
+ The Nal profile presented in Figure 1. has never been seen before in XX Oph., The NaI profile presented in Figure 1 has never been seen before in XX Oph.
+ While the sharp emissions at —37 kms + and the interstellar components at —11 km J| were already. known. the broad components at —365 km | are entirely new T'Darasov 2006. Goswami et 22001. Merrill 1951).," While the sharp emissions at $-$ 37 km $^{-1}$ and the interstellar components at $-$ 11 km $^{-1}$ were already known, the broad components at $-$ 365 km $^{-1}$ are entirely new Tarasov 2006, Goswami et 2001, Merrill 1951)."
+ There is a large uncertainty about (he reddening affecüng NX Oph., There is a large uncertainty about the reddening affecting XX Oph.
+ Lockwood et al. (, Lockwood et al. (
+1975) derived By 221.3. while Evans et al. (,"1975) derived $E_{B-V}$$\approx$ 1.3, while Evans et al. ("
+"1993) preferred. Lp,220.5.",1993) preferred $E_{B-V}$$\approx$ 0.5.
+ The IXKI 7699 iinterstellar line at —11 km + has an equivalent width of 0.186 40.007 oon our spectrum., The KI 7699 interstellar line at $-$ 11 km $^{-1}$ has an equivalent width of 0.186 $\pm$ 0.007 on our spectrum.
+ Using the calibration by Munari and Zwitter (1997). this translates into a reddening £g (420.73 40.03.," Using the calibration by Munari and Zwitter (1997), this translates into a reddening $E_{B-V}$ =0.73 $\pm$ 0.03."
+ This value agrees with the intensity of the Dilfuse Interstellar Bands at 5780. 5191. 5850. 6196. 6203 and 6614 {that are visible on our spectrum and with the core-saturated interstellar components of Nal in Figure 1.," This value agrees with the intensity of the Diffuse Interstellar Bands at 5780, 5797, 5850, 6196, 6203 and 6614 that are visible on our spectrum and with the core-saturated interstellar components of NaI in Figure 1."
+ The top panel of Figure 2 shows the spectrum of NX Oph over the wavelength range covered by the ongoing RAVE survey of the southern sky (Steinmetz et al., The top panel of Figure 2 shows the spectrum of XX Oph over the wavelength range covered by the ongoing RAVE survey of the southern sky (Steinmetz et al.
+ 2006. Zwitter el al.," 2006, Zwitter et al."
+ 2008) and the coming ESA’s GALA space mission (Munari 2003)., 2008) and the coming ESA's GAIA space mission (Munari 2003).
+ The same wavelength interval is also tentatively base-lined for a later phase of the LAMOST survey., The same wavelength interval is also tentatively base-lined for a later phase of the LAMOST $\footnote{http://www.lamost.org/website/en}$.
+ The bottom panel of Figure 2 shows the spectrum of XX Oph over the bluest (47084893 A)) of the four wavelength ranges to be covered by the forthcoming HIERMES all-sky survey’ (Freeman et 22010)., The bottom panel of Figure 2 shows the spectrum of XX Oph over the bluest (4708–4893 ) of the four wavelength ranges to be covered by the forthcoming HERMES all-sky $\footnote{http://www.aao.gov.au/AAO/HERMES/}$ (Freeman et 2010).
+ The other three ranges are 56495873. 6481.6739. and 75907890 A.. all covered bv the present. atlas.," The other three ranges are 5649–5873, 6481–6739, and 7590–7890 , all covered by the present atlas."
+ We are grateful to Federico Boschi aud Paola Marrese for securing some spectra, We are grateful to Federico Boschi and Paola Marrese for securing some spectra
+To describe the cluster density distribution we use the simple mass profile where is given by an ideutity similar to eq. (,To describe the cluster density distribution we use the simple mass profile where$m$ is given by an identity similar to eq. (
+13). with ay2>es: in the case of the spherical cluster we assnne ay=as.,"13), with $\acu\geq\acd\geq\act$; in the case of the spherical cluster we assume $\acu =\acd =\act$."
+ In the two following Sections we will compute the CTF components associated with eq. (, In the two following Sections we will compute the CTF components associated with eq. (
+19).,19).
+ Iu Appendix D it is shown that the tidal field compoucuts at the center of à non-uugular homeoidal distribution (such as that iu eq. [, In Appendix B it is shown that the tidal field components at the center of a non-singular homeoidal distribution (such as that in eq. [
+19|]. are given by Note that the quantities d; do not depend on the specific densitv profile. and that wy<aes for ay2a3. thus fulfilling the conditions for stable equilibriun in eq. (,"19]), are given by Note that the quantities $\wci$ do not depend on the specific density profile, and that $\wcu\leq\wcd\leq\wct$ for $\acu\geq\acd\geq\act$, thus fulfilling the conditions for stable equilibrium in eq. ("
+1).,1).
+" In analogy with definitions (15). we introduce the two cluster cllipticities. ως—1p and osea,—1p. πας pxosxld. and the expansion of the ¢uantities «e; for inall flatteniugs is given in Appendix C. Tn order to determine he galaxy oscillation frequencies we need. according to eqs. ("," In analogy with definitions (15), we introduce the two cluster ellipticities, $\acd/\acu\equiv 1-\mu$ and $\act/\acu\equiv 1-\nu$, with $\mu\leq\nu\leq 1$, and the expansion of the quantities $\wci$ for small flattenings is given in Appendix C. In order to determine the galaxy oscillation frequencies we need, according to eqs. ("
+3) and (20). to have a realistic estimate of ρω.,"3) and (20), to have a realistic estimate of $\rhocz$."
+ Unfortunately this quantity is not well measured iu real clusters. aud for its determination we Use the virial theorem. Mos5=U ⇁⇁(0 is the cluster gravitational potential energy. aud o- is the virial velocity dispersion. that we assume to be estimated by the observed velocitydispersion of galaxies iu the cluster).," Unfortunately this quantity is not well measured in real clusters, and for its determination we use the virial theorem, $\Mc\sigvd = -U$ $U$ is the cluster gravitational potential energy and $\sigvd$ is the virial velocity dispersion, that we assume to be estimated by the observed velocity dispersion of galaxies in the cluster)."
+" The explicit caleulatious of the factor 22Gp ix preseuted in Appendix D. We now compare the galactic oscillation periods P.= ὁπίω,, Py=2rfey. aud Pp=πω, with the characteristic stellar orbital times 2.1, iu galaxies: for siuplicity we eive here the expansions for αμα. cluster and galactic flatteuiues."," The explicit calculations of the factor $2\pi G\rhocz$ is presented in Appendix B. We now compare the galactic oscillation periods $\Pphi
+=2\pi/\omegphi$ , $\Pthe =2\pi/\omegthe$, and $\Ppsi =2\pi/\omegpsi$, with the characteristic stellar orbital times $\Porb$ in galaxies: for simplicity we give here the expansions for small cluster and galactic flattenings."
+ From eqs. (, From eqs. (
+3). (11). (20). aud uxiug eqs. (,"3), (14), (20), and using eqs. ("
+CLS)-(C.12) we finally obtains to the leading order in the fattcenines where we normalized ay to 1000 Lkin/s and e4 to 250 kpc.,C.8)-(C.12) we finally obtains to the leading order in the flattenings where we normalized $\sigv$ to 1000 km/s and $\acu$ to 250 kpc.
+ So. a comparison of quantities| above with οςB (," So, a comparison of quantities above with eq. ("
+18) shows that in the outer halo of giant Es the stellar orbital times cau be of the same order of magnitude as the oscillatory periods of the galaxies themselves.,18) shows that in the outer halo of giant Es the stellar orbital times can be of the same order of magnitude as the oscillatory periods of the galaxies themselves.
+" For example. in a relatively πα(Heruquist) galaxy of Magy=0.1 aud ayy=1. Paszm1 Cur at iiz10 (1.6.. atc! 5R,.). whilefor a galaxy with M44= Landay,=3the same orbital time is atf in27 (1.6.. at 2 3.5R.)."," For example, in a relatively small (Hernquist) galaxy of $\Mgn =0.1$ and $\alphaun =1$, $\Porb\simeq 1$ Gyr at $m\simeq 10$ (i.e., at $\simeq
+5\Reff$ ), whilefor a galaxy with $\Mgn =1$ and $\alphaun =3$the same orbital time is at $m\simeq 7$ (i.e., at $\simeq 3.5\Reff$ )."
+ Iu the case of spherical cluster. the CTF tensor im the reference system C. introduced in Sect;," In the case of spherical cluster, the CTF tensor in the reference system $\Cc$ introduced in Sect."
+" 2.2 is where qr)=petr)pe(r) aud pe=BAL)Lx? (sec, ce. CDIL CO9S)."," 2.2 is where $q(r) \equiv \rhoc (r)/\rhocm (r)$ and $\rhocm\equiv 3\Mc
+(r)/4\pi r^3$ (see, e.g., CD94, CG98)."
+ From eq. (, From eq. (
+19) we obtain wherePOLrfay.,19) we obtain where$\rt\equiv r/\acu$.
+ In analogy with Sect., In analogy with Sect.
+ 3.1. we now compare the galaxy oscillation periods with the mean stellar orbital times. aud for simplicity we restrict to the special case of the galaxy oscillating around. the €;-axis only. Le... we fix y=Ji and ο=0 in eq. (," 3.1, we now compare the galaxy oscillation periods with the mean stellar orbital times, and for simplicity we restrict to the special case of the galaxy oscillating around the $\xi_3$ -axis only, i.e., we fix $\varphi =0$ and $\vartheta =0$ in eq. ("
+10).,10).
+ Thus. from the results of Sect.," Thus, from the results of Sect."
+ 2.2. and from eq. (," 2.2, and from eq. ("
+21) where we used theAe last ofgi) eqs. (,24) where we used the last of eqs. (
+3) aud the expansion for stuall fattemimes of the coefficient involving the galaxy nluoluents of inertia.,3) and the expansion for small flattenings of the coefficient involving the galaxy moments of inertia.
+ Iu Fig., In Fig.
+ l we plot P. as a function of + for the two spherical cluster models used iu the siauulatious described in Sect., 1 we plot $\Ppsi$ as a function of $r$ for the two spherical cluster models used in the simulations described in Sect.
+ 5.2. in the cases of au E2 au and El ealaxv: note that P. decreases at mereasiug ealaxy flattening. while its radial trend is non monotonic. with a muinnmuua at pomay. the cluster core radius.," 5.2, in the cases of an E2 and and E4 galaxy; note that $\Ppsi$ decreases at increasing galaxy flattening, while its radial trend is non monotonic, with a minimum at $r\simeq a_1$, the cluster core radius."
+ The peculiarity of this location has been already: noted. (6.9.. sec CD91). and its impact on orbital evolution will be evideut iu Sect.," The peculiarity of this location has been already noted (e.g., see CD94), and its impact on orbital evolution will be evident in Sect."
+ 5.2, 5.2.
+ Iu the πιοσα simulations we use as time normalization the Πο]dle time fj=15 Civis. as Hass normalization the ealaxy iiass M. and as leneth normalization o1. the homeoidal scale-leneth corresponding to the galaxy major axis.," In the numerical simulations we use as time normalization the Hubble time $\tH =15$ Gyrs, as mass normalization the galaxy mass $\Mg$, and as length normalization $\alphau$, the homeoidal scale-length corresponding to the galaxy major axis."
+ Accordingly. eq. C," Accordingly, eq. ("
+I) in dimensionless form reads where 7=t/fyp aud o.=(GAL.fayjos: note that from eq. (,4) in dimensionless form reads where $\tau\equiv t/\tH$ and $\phig\equiv (G\Mg/\alphau)\tphig$; note that from eq. (
+3) uw;=wea; Where ως—28256p. and so from eqs. (,"3) $\omegi =\omegac\tomegi$ where $\omegac\equiv \sqrt
+{2\pi G\rhocz}$, and so from eqs. ("
+"2) aud (5) OQ=x, OQ.",2) and (5) $\Omegv=\omegac\tOmegv$ .
+ Iu addition. consisteutlv with the adopted “simall oscillatious approximation”. the angular velocity is used in linearized fortu Le. notein eq. (," In addition, consistently with the adopted “small oscillations approximation”, the angular velocity is used in linearized form i.e., in eq. ("
+5) we retained only linear terms in q. νε. CN that we neglect the direct effect of the cluster gravitational field on stellar orbits (cf.,"5) we retained only linear terms in $\phiM$ , $\theM$ , $\psiM$ ; note that we neglect the direct effect of the cluster gravitational field on stellar orbits (cf."
+ with eq. |, with eq. [
+6]).,6]).
+target gas giant.,target gas giant.
+ If this οποιον is efficiently converted to thermal οποιον throughout the merged core. temperature in (he core would increase by where 7; is the core temperature prior to the impact.," If this energy is efficiently converted to thermal energy throughout the merged core, temperature in the core would increase by where $T_i$ is the core temperature prior to the impact."
+ In the models La. ATne>5x10! K which is most likely to be larger than that (Tip) for heavy. elements to fully evaporate into a eas phase.," In the models La, $\Delta T_{\rm core} > 5\times 10^4$ K which is most likely to be larger than that $T_{\rm evap}$ ) for heavy elements to fully evaporate into a gas phase."
+ In this limit. the vapors of heavy elements would thoroughly mix with the nearby hydrogen gas in the envelope.," In this limit, the vapors of heavy elements would thoroughly mix with the nearby hydrogen gas in the envelope."
+ In the SPI models where (he core and the embryos are approximated as fluids with Tillotson equation of state For iron and basalt., In the SPH models where the core and the embryos are approximated as fluids with Tillotson \citep{Tillotson1962} equation of state for iron and basalt.
+ At lower temperature. the mixture of composite elements undergoes phase separation.," At lower temperature, the mixture of composite elements undergoes phase separation."
+ Although the phase separation diagram is well established [or H-IIe mixture 1976).. its analogs for II-Fe or H-Dasalt mixture are not available.," Although the phase separation diagram is well established for H-He mixture \citep{Stevenson1976}, its analogs for H-Fe or H-Basalt mixture are not available."
+ Nonetheless. we expect the heavy elements in the eas phase (o separate [rom hydrogen into a two-phase medium as the temperature decreases below a critical value 754.," Nonetheless, we expect the heavy elements in the gas phase to separate from hydrogen into a two-phase medium as the temperature decreases below a critical value $T_{2p}$."
+ During the phase separation. there is an exchange between Gibbs [ree energy and the internal οποιον 7; (Stevenson1976).," During the phase separation, there is an exchange between Gibb's free energy and the internal energy $E_i$ \citep{Stevenson1976}."
+. After phase separation. density contrast between the two phase increases with additional reduction in temperature.," After phase separation, density contrast between the two phase increases with additional reduction in temperature."
+ If the heavy elements can assemble liquid drops centered around (he nucleation sites and grow to sufficiently large sizes. (μον would precipitate and settle to the core (Nucleation is a poorly understood. critical phenomenon ancl it can proceed over a considerable (ime scale).," If the heavy elements can assemble liquid drops centered around the nucleation sites and grow to sufficiently large sizes, they would precipitate and settle to the core (Nucleation is a poorly understood critical phenomenon and it can proceed over a considerable time scale)."
+ Further reduction in planets internal temperature below some threshold μα would lead to the condensation of the heavy elements (IIubbard.1934)., Further reduction in planet's internal temperature below some threshold $T_{\rm cond}$ would lead to the condensation of the heavy elements \citep{Hubbard1984}.
+. We consider that for less massive CV;S a few AZ) impactors. especially for oblique collisions. (he impactors may disintegrate before reaching the core.," We consider that for less massive $M_I \lesssim$ a few $M_\oplus$ ) impactors, especially for oblique collisions, the impactors may disintegrate before reaching the core."
+" The distribution of heavy elemental deposition dMq,/dr depends on the impact angle. internal density inside the gas ejant. and (he mass of the impactor,"," The distribution of heavy elemental deposition $d M_{\rm
+dep} / dr$ depends on the impact angle, internal density inside the gas giant, and the mass of the impactor."
+ Without the loss of generality. we adopt a simple algebraic expression such Chat the rate of density added at each radius is where limp and 7445 represent the epoch of the impact and the duration of disintegration.," Without the loss of generality, we adopt a simple algebraic expression such that the rate of density added at each radius is where $t_{\rm imp}$ and $\tau_{\rm dep}$ represent the epoch of the impact and the duration of disintegration,"
+one). and it is still reasonable as for surface coverage (as shown iu Fig. 5)).,"one), and it is still reasonable as for surface coverage (as shown in Fig. \ref{fig:photo}) )."
+ The modoeliug has provided us with detailed information ou the heating deposition. both as for the spatial distribution aud for the temporal evolution.," The modeling has provided us with detailed information on the heating deposition, both as for the spatial distribution and for the temporal evolution."
+ A heat pulse deposited in the coronal part of loop A seenis unable to fit the sharp peak of the light curve., A heat pulse deposited in the coronal part of loop A seems unable to fit the sharp peak of the light curve.
+ A heating deposited at the loop footpoints is msteac more successful., A heating deposited at the loop footpoints is instead more successful.
+ Ou the other haud a heating at the footpoiuts is unable o dive the observed slow late decay. which iustead seems ο require a coronal location.," On the other hand a heating at the footpoints is unable to drive the observed slow late decay, which instead seems to require a coronal location."
+ The preseuce of both a steep rise phase aud a slow ate decav therefore suseests that kids of heating depositious. oue at he footpoiuts aud the other iu corona. nust be at work.," The presence of both a steep rise phase and a slow late decay therefore suggests that kinds of heating depositions, one at the footpoints and the other in corona, must be at work."
+ We cau also infer the relative weight of the two heating coniponeuts., We can also infer the relative weight of the two heating components.
+ The footpoiut heating is more impulsive. 1.6. intense aud short-lasting (a few miu).," The footpoint heating is more impulsive, i.e. intense and short-lasting (a few min)."
+ The other heating coniponent is less intense ( 1/1) aud releases its energy over a much longer time scale (one hour)., The other heating component is less intense $\sim 1/4$ ) and releases its energy over a much longer time scale (one hour).
+ Such features seein to be traced also by the optical light curve (see Paper ID. which shows a sharp peak aud a slower decay starting at ~L/L of the maxiumun optical count rate.," Such features seem to be traced also by the optical light curve (see Paper II), which shows a sharp peak and a slower decay starting at $\sim 1/4$ of the maximum optical count rate."
+ It is interesting to note that: a) the two compoucuts contribute to the flare with comparable amouuts of energy (2.3.10° ere): b) a simular combination of these two heating components (2210°? ere each) in the loop arcade (loop B) is able to explain he second flare maxima.," It is interesting to note that: a) the two components contribute to the flare with comparable amounts of energy $2 - 3 \times 10^{32}$ erg); b) a similar combination of these two heating components $\approx
+10^{32}$ erg each) in the loop arcade (loop B) is able to explain the second flare maximum."
+ Although we do not exclude that refining the heating function of the arcade may further inuprove the fitting of the data. we have shown hat usine simply the same time parameter values of the heating function vields a satisfactory description of the flare.," Although we do not exclude that refining the heating function of the arcade may further improve the fitting of the data, we have shown that using simply the same time parameter values of the heating function yields a satisfactory description of the flare."
+ihen77.,then.
+51) ence. most of. the energy is. radiated. [ar. away [rom. the NS surface.," Hence, most of the energy is radiated far away from the NS surface."
+ One mav also note the strong dependence of dL/d(logD) on the NS spin and radius., One may also note the strong dependence of $dL/d(\log{D})$ on the NS spin and radius.
+ The (ransition radius. Eq. (3.1)).," The transition radius, Eq. \ref{d-tr}) ),"
+ scales as, scales as.
+" The total boundary laver luminosity is then estimated to be: mam,!.(53) which perfectly agrees with a simple energeticargument that mir,! (Medvedev&Naravan9001)."," The total boundary layer luminosity is then estimated to be: m m, which perfectly agrees with a simple energeticargument that $L_{BL}\sim\dot Mc^2/r\sim \dot M_{\rm Edd}\dot m/r_*\sim m \dot m r_*^{-1}$ \citep{MN01}."
+ Our self-similar solution is subsonic (otherwise a shock would form) with the Mach number of the infalling gas beingdi. ⊔⋯↴↥↕⋝∖⊽⋅∐⋟⊔∐↲↖⊂≟≀↧↪∖⊽↕⊳∖⊽⋝∖⊽∏∣↽≻⊳∖⇁∪↕∏≺∢∖∖↽∐≼↲↕⋅≼↲⊔∐↲∣↽≻∪∏∐≼⇂≀↧↴↕⋅∡∖⇁↥≀↧∶∖⇁≼↲↕⋅∐⋯↴∩∢↥∐↲⊳∖⇁⊔∐↲∣↽≻∏∐≶≀↧↴≺∢≺∢↕⋅≼↲∐∪∐∐∪∖∖↽ ⋜⋝∖∖⊽∐↕≺∢∐↕⊳∖⊽↙∣∣∣∣⋎↙∣∫∕⊽↦↽⋅⊳∖⊽∪⋅∣↽≻≼↲≺∢≀↧↴∏⋝∖⊽≼↲⊔∐↲≸↽↔↴≀↧↪∖⊽↕∐⊔∐↲∐∪↥⊳∖⊽≼↲⊔∐∐≸≟∐∪∖∖↽↕⋝∖⊽∐↕≸↽↔↴⊔⋡∖↽⊳∖⇁∏∣↽≻⊳∖⇁∪↕," Our self-similar solution is subsonic (otherwise a shock would form) with the Mach number of the infalling gas being, that is, if the gas is subsonic where the boundary layer matches the bulk accretion flow (which is so, because the gas in the hot settling flow is highly subsonic), it will remain subsonic all the way down to the stellar surface."
+∏≺⋮⋟⋅∐∖∖↽↕∐↕⋅≼↲∐↓≀↧↴↕∐ ⊳∖⇁∏∣↽≻⊳∖⇁∪∐↕≺∢≀↧↴∐⊔∐↲∖∖↽," Next, we estimate the effect of Comptonization."
+≀↧∶∖↽≼⇂∪∖∖⊽∐↥∪⊔∐↲⊳∖⇁∩↲∐≀↧↴↕⋅⊳∖⊽⋯⋅↓⋟≀↧↴≺∢≼↲⋅ ↳∖↡≼↲⇀↸↥⋅∖∖↽≼↲≼↲⊳∖⇁∐∐⋯∥↲⊔∐↲≼↲∐⋡≼↲≺∢↥∪↓⋟≼↕⊲∪↕∐↕↽≻↥∪∐↕∠≀↧↴∐∪∐⋅↴∏∐↲≼↲↥≼↲≺∢∏⋅∪∐⋝∖⊽≺∢≀↧↴⊔≼↲↕⋅↕∐≸↽↔↴∪↕↽≻∐≺∢≀↧↴↥≼⇂≼↲↕↽≻⊔↥↥⊳∖⊽ ⊤≺∖∖≃∕↗∣≖⇁≺∖∖∐⋅∖∖," The electron scattering optical depth is $\tau_{\rm es}\simeq\rho\kappa_{\rm es}D$, where $\kappa_{\rm es}=\sigma_T/m_p$ is the electron scattering opacity for ionized hydrogen."
+↽∐≼↲↕⋅≼↲∣↘⇁≺⋅⋝∖∶⊔⊻⊽∕∕∕∣∣∣↙⇂↥⊳∖⊽⊔∐↲≼↲↥≼↲≺∢∏⋅∪∐⊳∖⇁≺∢≀↧↴⊔≼↲↕⋅↕∐≸≟∪↕↽≻≀↧↴≺, We can use our self-solution for $\rho$ to calculate $\tau_{\rm es}$.
+"∢∐⋡∖↽↓≯∪↕⋅↕∪∐↕∠≼↲≼⇂∐∡∖⇁≼∐⋅∪≸≟≼↲∐⋅∖∖⊽≼↲ ≺∢≀↧↴∐∏⊳∖⊽≼↲⋯∐⋅⊳∖⇁≼↲∐⋟−⋝∖⊽∪↥∏∐∪∐↓⋟∪↕⋅∕↗↥∪≺∢≀↧↴↥≺∢↕∐≀↧↴∥↲⊤≺⋅⊓⋅↼≚∐≼↲↕⋅∐≀↧↴∐∖↽≼↲↥∡∖↽⋅∖∖↽≼↲↕∐≀↧∶∖↽↕⋅≼↲≺∢≀↧↴∐⊔⋯↴↥⊔∐↲≼⇂≼↲∐⊳∖⇁∐∡∖↽ matches that of the hot settling flow at the transition distance. and at in (his region the electron scattering optical depth (\ledvedev&Naravan 2001)..τιςcLO%as?r,+ (asstuming again that Dye< Ry). is S1/3 for typical parameters a=0.1. s=0.1 and r,= 3."," Alternatively, we may recall that the density matches that of the hot settling flow at the transition distance, and at in this region the electron scattering optical depth \citep{MN01}, ,$\tau_{\rm es}\simeq 10^3\alpha s^2 r_*^{-1}$ (assuming again that $D_{\rm tr}\ll R_*$ ), is $\lesssim1/3$ for typical parameters $\alpha=0.1$ $s=0.1$ and $r_*=3$ ."
+ Using, Using
+the combined Xa:w luminosity but we divide that value in half to generate the N-rav Iuninositv function.,the combined X-ray luminosity but we divide that value in half to generate the X-ray luminosity function.
+ Tn order to «mupare the derived X-ray luminosity function for the TTS in MDMI2 with other flux liuüted N-rav hpnuauimositv fuctions we used the ASURV Rev. 1.2 package (]sobe Feigelson 1990: LaValley et al., In order to compare the derived X-ray luminosity function for the TTS in MBM12 with other flux limited X-ray luminosity functions we used the ASURV Rev. 1.2 package (Isobe Feigelson 1990; LaValley et al.
+ 1992). which implements the ucthods preseutec in Feieclson Nelson (1985).," 1992), which implements the methods presented in Feigelson Nelson (1985)."
+ Although he currently known TTS nu AIDMI2 are a] Nav detected. tio Dunrosity functions of other. uore distaut. star forie regions include upper limuts.," Although the currently known TTS in MBM12 are all X-ray detected, the luminosity functions of other, more distant, star forming regions include upper limits."
+ The derived N-ray hinunosity fuuction is displayed in Fig., The derived X-ray luminosity function is displayed in Fig.
+ δ with fhe N-ray huuirosity function for the TTS in the L1195E cloud 1 Taurus which (like MDBMI2) was observed in a decp (33 kx)ROSAT PSPC Toed observation (Sroni Stron 199y," \ref{lumfunc}
+ with the X-ray luminosity function for the TTS in the L1495E cloud in Taurus which (like MBM12) was observed in a deep (33 ks) PSPC pointed observation (Strom Strom 1994)."
+ TheROSAT TOWed observation of LEI95E is ~ 20 times more sesitive than previous observatious with the satellite., The pointed observation of L1495E is $\sim$ 20 times more sensitive than previous observations with the satellite.
+ Strom Strom (1991) used εμας observation to show tiat the Nauav hinüiuostv of LTS extends to fainter nuimosities than were observec| withEiusteiu., Strom Strom (1994) used this observation to show that the X-ray luminosity of TTS extends to fainter luminosities than were observed with.
+ We have re-reduced the pointed observatiou of L1195E in a way analogous to that of AIBA12., We have re-reduced the pointed observation of L1495E in a way analogous to that of MBM12.
+ The X-ray huuinositv function we derive for L1195E (1) inclucess only he I& aud M spectral type TTS. (2)1ichdes 6 upper lanits. (3) asses an X-ray to optical flux conversion factor of 1.1. « Moye tom 7. aud (1) asstmes a distance o: 110 pe.," The X-ray luminosity function we derive for L1495E (1) includes only the K and M spectral type TTS, (2) includes 6 upper limits, (3) assumes an X-ray to optical flux conversion factor of 1.1 $\times$ $^{-11}$ erg $^{-1}$ $^{-2}$, and (4) assumes a distance of 140 pc."
+ The X-ray huninositv ‘tions in MDMI2 aud L1195E agree well: in MDMI2 ogLxnncalr 29.040.1 cre land logL.dan 28.7 ere tt in L1195E logeLywean 9+0.2 cre st and logelsqueian 28.9 ere st., The X-ray luminosity functions in MBM12 and L1495E agree well: in MBM12 the $L_{\rm x~mean}$ = $\pm$ 0.1 erg $^{-1}$ and $L_{\rm x~median}$ = 28.7 erg $^{-1}$; in L1495E $L_{\rm x~mean}$ = $\pm$ 0.2 erg $^{-1}$ and $L_{\rm x~median}$ = 28.9 erg $^{-1}$.
+ However. we uote hat the NDMI2 ταν numositv function has a lower Hel-huninosity liwit and a hieher low-huuinosity limit han the L1195E N-rav luminosity. function.," However, we note that the MBM12 X-ray luminosity function has a lower high-luminosity limit and a higher low-luminosity limit than the L1495E X-ray luminosity function."
+ Therefore. although the poiited observation of MDMI2 Is lore sensitive than the pointed observation LLLOSE (because AIDMAI2 is uch closer} our follow-up obscrvatious of the TTS in MDMI2 nay be incomplete for sources fainter than V — 15.5 nag.," Therefore, although the pointed observation of MBM12 is more sensitive than the pointed observation L1495E (because MBM12 is much closer) our follow-up observations of the TTS in MBM12 may be incomplete for sources fainter than $V$ $\sim$ 15.5 mag."
+" Iu addition. since we know that one of our N-rav sotτοῦ», SIS. is deteced but below our lreshold for follow-up observations. there may be other. ater, N-rav οτίing TTS in MBA[12 with spectral vpes later than ~ A2 (ic.. the spectra type of S18) that will be discovered in nore sensitive folow-up observations."," In addition, since we know that one of our X-ray sources, S18, is detected but below our threshold for follow-up observations, there may be other, fainter, X-ray emitting TTS in MBM12 with spectral types later than $\sim$ M2 (i.e., the spectral type of S18) that will be discovered in more sensitive follow-up observations."
+ The discrepancy at he high huninositsond of the X-ray mnuimositv fuuctiou nav also be explaa.τος if the distauce o the TTS in MDMI2 is lavecr than 65 pe., The discrepancy at the high luminosity end of the X-ray luminosity function may also be explained if the distance to the TTS in MBM12 is larger than 65 pc.
+ Although anu mereased clistance is allowed by he' yecent. results it should be confirmed with furher observations., Although an increased distance is allowed by the recent results it should be confirmed with further observations.
+ Although MDMI2 is not a prolific star-forming cloud when compared to nearby giant molecular clouds it is he nearest star-forming cloud to the sun aud offers a unique opportuv to study the star-formation process within a nolecular cloud :uo hieh seusitiviv., Although MBM12 is not a prolific star-forming cloud when compared to nearby giant molecular clouds it is the nearest star-forming cloud to the sun and offers a unique opportunity to study the star-formation process within a molecular cloud at high sensitivity.
+ We have xeseuted followp observalous of Xa‘av stars identified in the region of the MDMAI2 complex., We have presented follow-up observations of X-ray stars identified in the region of the MBM12 complex.
+ These observations iive doubled the nunber of confirmed TTS i1 this region., These observations have doubled the number of confirmed TTS in this region.
+ Since theROSAT PSPC poiited observation of the ceutral roeion oft 1ο clo Was sensiive enough to detect all of the xeviouslv known TTS in the cloud. we believe the list of ος ρα id 5 WTTS presented in Table { to © a nearlv complee census of the TTS in MDMI2 for sρουτα] types earlier hau  M2.," Since the PSPC pointed observation of the central region of the cloud was sensitive enough to detect all of the previously known TTS in the cloud, we believe the list of 5 CTTS and 3 WTTS presented in Table \ref{eqw} to be a nearly complete census of the TTS in MBM12 for spectral types earlier than $\sim$ M2."
+" Assumi1 a luca mass of ~ 066 M for the S currently kuown TTS iu MDMI?2 aid a cloud mass of 30.200 ML, (Pound et al.", Assuming a mean mass of $\sim$ 0.6 $_{\odot}$ for the 8 currently known TTS in MBM12 and a cloud mass of 30–200 $_{\odot}$ (Pound et al.
+ 199k Ziumeriamn Cueerechts 1990) he star-formation efficiency of NDMI2 js , 1990; Zimmermann Ungerechts 1990) the star-formation efficiency of MBM12 is $\sim$.
+"Since the currently known TTS »opulation in MDMI?2 is incouplete ouly for the lower mass ob;jects, unless there are a huge mmuber of these objecs vot to be cüscovered i the coud. this estimate of the star-formation eficicney will not change siguificautlv."," Since the currently known TTS population in MBM12 is incomplete only for the lower mass objects, unless there are a huge number of these objects yet to be discovered in the cloud, this estimate of the star-formation efficiency will not change significantly."
+ Althoteh there is still a large uncertaintv in the mass of the cloud the estimated star-fornation cficiencics are consisten with that expected from clouds with masses on tre order of 100 AL. (Ehuecercen Efremov 1997)., Although there is still a large uncertainty in the mass of the cloud the estimated star-formation efficiencies are consistent with that expected from clouds with masses on the order of 100 $_{\odot}$ (Elmegreen Efremov 1997).
+ By compariie the PAreneths of the Πα emission and A6708 aabsorption lines of the TTS in MDMI2 with hose found oei other voung clusters. we place au upper liuüt on ic aee of the stars in MDMI2 ~ 10 My.," By comparing the strengths of the $\alpha$ emission and $\lambda$ 6708 absorption lines of the TTS in MBM12 with those found in other young clusters, we place an upper limit on the age of the stars in MBM12 $\sim$ 10 Myr."
+ By comparing the Nav luminosity fiction of ic TTS in MDMI2 with that o© the. TTS iu L1195E we predict that there are more vouug. low-niass. stars to TE discovered in MDMI2 aud the assuued distance to ic cloud παν have to be increased.," By comparing the X-ray luminosity function of the TTS in MBM12 with that of the TTS in L1495E we predict that there are more young, low-mass, stars to be discovered in MBM12 and the assumed distance to the cloud may have to be increased."
+ Althotwh this prediction agrees with the recently revised distance estinate to the cloud (~ 65435 pe) based oj results of the satellite. it should be coufirines with future ooervatious.," Although this prediction agrees with the recently revised distance estimate to the cloud $\sim$ $65\pm35$ pc) based on results of the satellite, it should be confirmed with future observations."
+ We have also identified a reddened C9 star behind the cloud with A.  8.18.9 mae., We have also identified a reddened G9 star behind the cloud with $A_{\rm v}$ $\sim$ 8.4–8.9 mag.
+ Therefore. here are at least two lines of sight through the cloud that show larger extinctious (el. > 5 mag) than previously hought for this cloud.," Therefore, there are at least two lines of sight through the cloud that show larger extinctions $A_{\rm v}$ $>$ 5 mag) than previously thought for this cloud."
+ This higher extinction explains why MDMIT2 is, This higher extinction explains why MBM12 is
+epochs are needed. in which the candidate is imaged together with TTel.,"epochs are needed, in which the candidate is imaged together with Tel."
+ Therefore. we searched first in the ESO data archive for additional NACO observations of TTel. and found a data-set in which the candidate is clearly detected.," Therefore, we searched first in the ESO data archive for additional NACO observations of Tel, and found a data-set in which the candidate is clearly detected."
+" The public archival data were taken in June 2007 with NACO's S27 optics through the K,-band filter in combination with NACO’s neutral density filter NDsnpow (transmission of about To determine the accurate astrometry of the candidate relative to TTel. images of astrometric calibrators have to be taken in each observing run."," The public archival data were taken in June 2007 with NACO's S27 optics through the $\rm K_{s}$ -band filter in combination with NACO's neutral density filter $\rm ND_{Short}$ (transmission of about To determine the accurate astrometry of the candidate relative to Tel, images of astrometric calibrators have to be taken in each observing run."
+ Therefore. in September 2009. we observed the globular cluster TTuc for which precise HST astrometry is available for several of its members.," Therefore, in September 2009, we observed the globular cluster Tuc for which precise HST astrometry is available for several of its members."
+" We determined the pixel scale PS and the position angle PA of NACO’s S13 optics. which is listed in candidatewecanreftable, strocaltogetherwiththeastrometriccalibrationo f thear Inee"," We determined the pixel scale $PS$ and the position angle $PA$ of NACO's S13 optics, which is listed in \\ref{table_astrocal} together with the astrometric calibration of the archival data, taken from \cite{chauvin2010}."
+ With the given astrometric calibration we could then determine the relative astrometry of the companion-candidate. Le. Its angular separation (sep) and position angle (PA) to TTel. in all available observing epochs.," With the given astrometric calibration we could then determine the relative astrometry of the companion-candidate, i.e. its angular separation $sep$ ) and position angle $PA$ ) to Tel, in all available observing epochs."
+ Due to the small separation of the candidate from the much brighter star. the PSF of TTel was always removed in all images with spatial filtering. to determine the relative astrometry of the candidate.," Due to the small separation of the candidate from the much brighter star, the PSF of Tel was always removed in all images with spatial filtering, to determine the relative astrometry of the candidate."
+ All astrometric measurements are summarized in reftableeppacandareillustratedinFig. refseppa.., All astrometric measurements are summarized in \\ref{table_seppa} and are illustrated in \\ref{seppa}.
+ The comparison of our NACO data from September 2009 with the archival data from June 2007 already proofs that the detected companion-candidate is clearly not a non-moving background source., The comparison of our NACO data from September 2009 with the archival data from June 2007 already proofs that the detected companion-candidate is clearly not a non-moving background source.
+ The expected separations and position angles for such an object can be derived (see reftable;eppa. .andFig. ref seppayfromtherelativeastrometryo Fthecandidate. thegivene poc," The expected separations and position angles for such an object can be derived (see \\ref{table_seppa}, and \\ref{seppa}) ) from the relative astrometry of the candidate, the given epoch differences, as well as the proper and parallactic motion of Tel, well known from Hipparcos."
+"hdi, (PA=-0.5+ 0.3°/yr). while a decrease of about 27° is expected for a non-moving background source."," We find that the position angle of the candidate only slightly decreases between both observing epochs by $1.19\pm0.62$ $^{\circ}$ $\dot{PA} = -0.5\pm0.3$ $^{\circ}$ /yr), while a decrease of about $^{\circ}$ is expected for a non-moving background source."
+ Hence. reject the background hypothesis for the detected already from its position angle measurements (on a chis(ee quitan?39 ιτ].," Hence, we can reject the background hypothesis for the detected candidate already from its position angle measurements (on a significance level $> 39\,\sigma$ )."
+ While the change in the positior angle is relatively small. the separation. of the candidate significantly increases. between both observing epochs by 82.3€2.0 mmas (sep=35.8+1.3 mmas/yr).," While the change in the position angle is relatively small, the separation of the candidate significantly increases between both observing epochs by $82.3\pm2.9$ mas $\dot{sep} = 35.8\pm1.3$ mas/yr)."
+ However. evel this high relative motion can be explained with orbital motior around TTel (see section 4 for discussion).," However, even this high relative motion can be explained with orbital motion around Tel (see section 4 for discussion)."
+ Therefore. we conclude that the detected candidate ts a real companion of TTel. and we will denote it as BB. from now on.," Therefore, we conclude that the detected candidate is a real companion of Tel, and we will denote it as B, from now on."
+ In order to follow the orbital motion of BB around its primary we observed the TTel system in May 2010 again with NACO in the Κ.-ραπά., In order to follow the orbital motion of B around its primary we observed the Tel system in May 2010 again with NACO in the $\rm K_{s}$ -band.
+ For the astrometric calibration images of the globular cluster 47TTue were taken. as well," For the astrometric calibration images of the globular cluster Tuc were taken, as well"
+The Fornax dwarf galaxy is the largest. of the dwarl spheroidal galaxies of the Milkv. Way.,The Fornax dwarf galaxy is the largest of the dwarf spheroidal galaxies of the Milky Way.
+" Hs integrated stellar luminosity is L,=(1.58c0.16)«10 L. (Mateo 1998)). but it has retained virtually no mass (< 500047.) in gas (Mateoetal. 19912). whieh makes pinpointing its radial velocity less accurate."," Its integrated stellar luminosity is $L_v=(1.58\pm0.16)\times 10^7\lsun$ \citealt{mateo98}) ), but it has retained virtually no mass $<5000\msun$ ) in gas \citealt{mateo91}) ), which makes pinpointing its radial velocity less accurate."
+ Reeardless. the best estimates (Mateoetal. 1991)) suggest a Leliocentric distance anc advancing radial velocity of 1384SApe and 53-43kms* respectively.," Regardless, the best estimates \citealt{mateo91}) ) suggest a Heliocentric distance and advancing radial velocity of $138\pm8~kpc$ and $53\pm3\kms$ respectively."
+" Fornax has a well measured surface density profile enabling the deduction of its Wine model parameters which define the number density of stars [rom the centre to the tical radius. A,=2.080.15 Ape."," Fornax has a well measured surface density profile enabling the deduction of its King model parameters which define the number density of stars from the centre to the tidal radius, $R_t=2.08\pm0.18$ $kpc$."
+" This is complemented: by the concentration parameter log,/F8)0.72 (Irwin&Llatzicimitriou 1995)).", This is complemented by the concentration parameter $\log_{10}(R_t/R_c)=0.72$ \citealt{irwinhatz}) ).
+" Stellar population synthesis models tell us that for the ages and metallicities of stars in Fornax. the ratio of mass to luminosity should. be of order. unity. (Ixroupa. 20013). however. the first puzzle of Fornax. as with the other dwarls. is that. relating the random motions of the stars to their implied. mass from their integrated luminosity leaves an incongruitv: the gravity GAL,(or is simply insullicient."," Stellar population synthesis models tell us that for the ages and metallicities of stars in Fornax, the ratio of mass to luminosity should be of order unity \citealt{kroupa01}) ), however, the first puzzle of Fornax, as with the other dwarfs, is that relating the random motions of the stars to their implied mass from their integrated luminosity leaves an incongruity: the gravity $GM_{\star}(r)r^{-2}$ is simply insufficient."
+" There is an ""acceleration celicit"""," There is an “acceleration deficit""."
+ Since we assume the measurements are correct. either eravitv is not. Newtonian. allowing the acceleration to be boosted according to some algorithm. or there is some dark matter (DM) which provides a similar service.," Since we assume the measurements are correct, either gravity is not Newtonian, allowing the acceleration to be boosted according to some algorithm, or there is some dark matter (DM) which provides a similar service."
+ What this means in detail is discussed in re[fsec:idm and & re[secimond.. nevertheless. the observations of Fornax have recentIy become so detailed (Walkeretal. 2007)) that more than two thousand member stars can be sorted into dillerent projected. radii bins giving the line of sight. (los). velocity dispersion. (VD) as a function. of projected. radius.," What this means in detail is discussed in \\ref{sec:dm} and \\ref{sec:mond}, nevertheless, the observations of Fornax have recently become so detailed \citealt{walker07}) ) that more than two thousand member stars can be sorted into different projected radii bins giving the line of sight (los) velocity dispersion (VD) as a function of projected radius."
+ This basically fixes the allowed mass density. luminous or clark. at all racii where stars exist.," This basically fixes the allowed mass density, luminous or dark, at all radii where stars exist."
+ The second. puzzle is far more subtle., The second puzzle is far more subtle.
+ In. orbit. there are at least 5 well resolved. globular clusters (GC's). with important parameters given in Table 1.," In orbit, there are at least 5 well resolved globular clusters (GCs), with important parameters given in Table 1."
+ The GC masses range between 0.37.107A7. and 3.63.107AZ. . the projected," The GC masses range between $0.37\times 10^5\msun$ and $3.63\times 10^5\msun$ , the projected"
+"107! 10 Fatatypical Whiteetal.1995 M. ~2x10"" !. ",$^{34}$ $^{38}$ $^{-1}$ \citealt{WNP95} $_\odot$ $\sim 2 \times 10^{38}$ $^{-1}$ 
+in cvele 22. the primary cireulation cell flowed poleward only to about 607. Iatitude. (hus making a shorter path for the magnetic Πας transport via (he convevor bell ancl resulting in aevele duration of 10.5 vears (see the left [rame of Figure 1).,"in cycle 22, the primary circulation cell flowed poleward only to about $60^{\circ}$ latitude, thus making a shorter path for the magnetic flux transport via the conveyor belt and resulting in a cycle duration of $\sim 10.5$ years (see the left frame of Figure 1)."
+ On the other hand. in evele 23. (he poleward surface flow went all the way to the poles (see the middle frame in Figure 1). leading to a longer path via the convevor-belt evele of ~12.5 vears.," On the other hand, in cycle 23, the poleward surface flow went all the way to the poles (see the middle frame in Figure 1), leading to a longer path via the conveyor-belt cycle of $\sim 12.5$ years."
+ In (he past. some flux-transport dynamo caleulations (Bonannoetal2005) and surface-transport calculations (Jiangetal2003). have dealt with possible multi-cell meridional flow scenarios. in the context of understanding the role of these flows in solar evcele features.," In the past, some flux-transport dynamo calculations \citep{bebr05} and surface-transport calculations \citep{jcss08} have dealt with possible multi-cell meridional flow scenarios, in the context of understanding the role of these flows in solar cycle features."
+ We now see (hat these studies are not just the merely plaving with models: such scenarios could happen in realitv., We now see that these studies are not just the merely playing with models; such scenarios could happen in reality.
+ The change in (he surface poleward flow-patlern its reversing around 607 as it did in evele 22 and maintaining poleward flow all the way to the pole as in evcle 23 have significantly impacted the duration of evele 23 ancl the length of the minimum that Iollowed it. compared to that in cvele 22.," The change in the surface poleward flow-pattern – its reversing around $60^{\circ}$ as it did in cycle 22 and maintaining poleward flow all the way to the pole as in cycle 23 – have significantly impacted the duration of cycle 23 and the length of the minimum that followed it, compared to that in cycle 22."
+ The plasma velocity can also be determined by helioseismic analvsis. either [vom ring ciagrams or from time-distance diagrams (Gilesetal1997:Haber2002:Zhao&Ixoso-vichey2004:Gonzalez-lIernandezetal2008:Gizon.Birch&Spruit 2010).," The plasma velocity can also be determined by helioseismic analysis, either from ring diagrams or from time-distance diagrams \citep{gdsb97,
+hetal02,zk04,gkhhk08,gbs10}."
+. Alternatively. features seen on (he images such as magnetic structures and supereranule cells. can be tracked with cross-correlation analvsis to. vield a drift velocity for that leature (Nomi.Howard&Harvey1993:Snodgrass&Dailey1996:Svandaetal2006.2007. 2008).," Alternatively, features seen on the images such as magnetic structures and supergranule cells, can be tracked with cross-correlation analysis to yield a drift velocity for that feature \citep{khh93,sd96,sks06,szk07,sksb08}."
+. Ulrich(2010) reanalyzed (he Mount Wilson surface Doppler data lor evcles 22 and 923. from 1986 through 2009.," \citet{u10} reanalyzed the Mount Wilson surface Doppler data for cycles 22 and 23, from 1986 through 2009."
+ Ulrich(2010) computed meridional flow profiles up to at least 80° latitude. found smooth evolution of the signal from one vear to the next. and confirmed the dillerence in hieh latitude flow patterns between cycles 22 and 23. in particular. the existence of a second reversed cell poleward of about 607 during most of evele 22. and a single cell with poleward surface flow all the wav to the poles during the major part of evcle 23.," \citet{u10} computed meridional flow profiles up to at least $80^{\circ}$ latitude, found smooth evolution of the signal from one year to the next, and confirmed the difference in high latitude flow patterns between cycles 22 and 23, in particular, the existence of a second reversed cell poleward of about $60^{\circ}$ during most of cycle 22, and a single cell with poleward surface flow all the way to the poles during the major part of cycle 23."
+ Differential rotation throughout the solar convection zone is fairly well known from helioseisnuc measurements (Thompson2004)., Differential rotation throughout the solar convection zone is fairly well known from helioseismic measurements \citep{t04}.
+.. It is difficult to measure al hieh latitudes because of foreshortening and other effects., It is difficult to measure at high latitudes because of foreshortening and other effects.
+ Most methods measure the linear rotational velocitv rather than (the angular rotation rate. so il is particularly difficult to calculate the angular measure with (he short moment-arm near the poles.," Most methods measure the linear rotational velocity rather than the angular rotation rate, so it is particularly difficult to calculate the angular measure with the short moment-arm near the poles."
+ It appears that on average the angular measure of differential rotation al the surface declines monotonically to the poles (Beck2000).. though the existence of a polar vortex’ has been suggested by (theory 1979).," It appears that on average the angular measure of differential rotation at the surface declines monotonically to the poles \citep{b00}, though the existence of a 'polar vortex' has been suggested by theory \citep{g79}."
+. Although our focus on this paper is more on meridional flow at high Iatitudes than on differential rotation there. our model will caleulate both quantities. so we will need to compare results from the model with both flows.," Although our focus on this paper is more on meridional flow at high latitudes than on differential rotation there, our model will calculate both quantities, so we will need to compare results from the model with both flows."
+ Given the important consequences of a second. reversed meridional cell in high latitudes.," Given the important consequences of a second, reversed meridional cell in high latitudes,"
+disk (Gyc=5 kpe).,disk $R_{\rm LMC}=5$ kpc).
+ His clear from (his Fie., It is clear from this Fig.
+ 3 that if there are only a few hundreds Ηλος within the LAIC orbit. then the IIVCs are highly unlikely to be accreted onto the LAIC and consequently dilute the ISM within a timescale of well less than 10vr (which corresponds to ages of voung stellar populations with low |N/II] in the LMC).," 3 that if there are only a few hundreds HVCs within the LMC orbit, then the HVCs are highly unlikely to be accreted onto the LMC and consequently dilute the ISM within a timescale of well less than $10^8$ yr (which corresponds to ages of young stellar populations with low [N/H] in the LMC)."
+ By assuming a (vpical mass of the individual HINVC's Gry.) aid using the results shown in Fig., By assuming a typical mass of the individual HVCs $m_{\rm hvc}$ ) and using the results shown in Fig.
+ 3. we can diseuss (he possible accretion rate of the ΗΝ onto the LMC disk for a eiven sel of model parameters.," 3, we can discuss the possible accretion rate of the HVCs onto the LMC disk for a given set of model parameters."
+ Fig., Fig.
+ 4 shows that Mie is much less than e0.1M.vr| for almost all models with different m and Nice., 4 shows that ${\dot{M}}_{\rm HVC}$ is much less than $\sim 0.1 {\rm M}_{\odot} {\rm yr}^{-1}$ for almost all models with different $m_{\rm hvc}$ and $N_{\rm HVC}$ .
+ The minimum value of the required. Vie shown in Fig., The minimum value of the required $M_{\rm HVC}$ shown in Fig.
+ | is L8x10?M. for fy=I0*vr in different models with different. ey and ονμις., 1 is $1.8 \times 10^6 {\rm M}_{\odot}$ for $t_{\rm sf}=10^7$ yr in different models with different ${\epsilon}_{\rm sf}$ and ${\rm [N/H]}_{\rm HVC}$.
+" Therefore. at least 0.18 M. | is necessary to dilute the ISM of the LAC to the observed level for /,;=LO"" vr."," Therefore, at least 0.18 ${\rm M}_{\odot}$ $^{-1}$ is necessary to dilute the ISM of the LMC to the observed level for $t_{\rm sf}=10^7$ yr."
+ Ht should be stressed that the above 0.18 M. | is for e=1.0 (i.e.. star Formation elliciency): a realistic value of the required Mauve is likely to be significantly larger than 0.18 M. |l," It should be stressed that the above 0.18 ${\rm M}_{\odot}$ $^{-1}$ is for ${\epsilon}_{\rm sf}=1.0$ (i.e., star formation efficiency): a realistic value of the required ${\dot{M}}_{\rm HVC}$ is likely to be significantly larger than 0.18 ${\rm M}_{\odot}$ $^{-1}$."
+ The results shown in Fig., The results shown in Fig.
+" 4 suggest that only models with verv large (vpical masses of ΗΝος (ie. mp= 10M.) and large number of the HIVCS. (Nive>2500) can show Myve as high as the required. rate above (0.138 AL. vr "")."," 4 suggest that only models with very large typical masses of HVCs (i.e., $m_{\rm hvc} = 10^7 {\rm M}_{\odot}$ ) and large number of the HVCs $N_{\rm HVC}>2500$ ) can show ${\dot{M}}_{\rm HVC}$ as high as the required rate above (0.18 ${\rm M}_{\odot}$ $^{-1}$ )."
+) Although the required. typical nass is similar (ο the observed mass of Complex C (e.g.. Thom et al.," Although the required typical mass is similar to the observed mass of Complex C (e.g., Thom et al."
+ 2003). the required total number within the LMCs orbital radius already exceeds the total number of the IINVC's (~ 2000) observed by the HII Parkes All Sky Survey (IIIPASS: e.g. Putman et al.," 2008), the required total number within the LMC's orbital radius already exceeds the total number of the HVCs $\sim 2000$ ) observed by the HI Parkes All Sky Survey (HIPASS; e.g., Putman et al."
+ 2002): it should be noted that the observed one is for the IIVCs existing in the entire regions around the Galaxy whereas (he required one is only for those within ~75 kpc., 2002): it should be noted that the observed one is for the HVCs existing in the entire regions around the Galaxy whereas the required one is only for those within $\sim 75$ kpc.
+ These results imply that it is unlikely for the accretion of the Galactic IIVCs onto the LAIC disk to dilute the ISM., These results imply that it is unlikely for the accretion of the Galactic HVCs onto the LMC disk to dilute the ISM.
+ However Άλμνο could become large enough in a sporadic wav if Cae LMC can interact with groups of IIVC's with locally laree number densities., However ${\dot{M}}_{\rm HVC}$ could become large enough in a sporadic way if the LMC can interact with groups of HVCs with locally large number densities.
+ Although we have adopted a reasonable range of model parameters and thereby investigated (1) the accretion rate of the IIVCs onto the LMC and (ii) the possible total mass of (he IINC's within the outer Galactic halo.there could be some uncertainties in model parameters.," Although we have adopted a reasonable range of model parameters and thereby investigated (i) the accretion rate of the HVCs onto the LMC and (ii) the possible total mass of the HVCs within the outer Galactic halo,there could be some uncertainties in model parameters."
+ Thus we here discuss how (he present results depend on these model parameters., Thus we here discuss how the present results depend on these model parameters.
+which averages towards zero as the integration length is increased.,which averages towards zero as the integration length is increased.
+" Similarly, if two references are used to create two independent la cancellers for the two main signal copies, this gives the MKMKIb canceller shown in figure 4,, and the mean residual output power becomes Equations (14)) and (15)) show that while R,»z.2,) should towards zero, £;,;;, has a definite limit due to keepthe RFI integratingsignal that remains after cancelling."," Similarly, if two references are used to create two independent MK1a cancellers for the two main signal copies, this gives the MK1b canceller shown in figure \ref{indepRX MK1}, and the mean residual output power becomes Equations \ref{mk2 indepRX}) ) and \ref{mk1 indepRX}) ) show that while $R_{\it{mk2b}}$ should keep integrating towards zero, $R_{\it{mk1b}}$ has a definite limit due to the RFI signal that remains after cancelling."
+" However, one must be aware that in situations where the reference INR is very small the MK2 canceller does not turn itself off, and there are practical implementation issues that need to be addressed (essentially, one may need to force the canceller to turn off)."," However, one must be aware that in situations where the reference INR is very small the MK2 canceller does not turn itself off, and there are practical implementation issues that need to be addressed (essentially, one may need to force the canceller to turn off)."
+ This is discussed in section 6..," This is discussed in section \ref{INSTABILITIES IN THE
+DUAL REFERENCE ALGORITHMS}."
+ In this section the residual equations given throughout section 2 are demonstrated., In this section the residual power equations given throughout section \ref{ADAPTIVE CANCELLERS} are demonstrated.
+"power Note that while the plots for the dual-reference cancellers show the residual power as the reference interference-to-noise power ratio approaches arbitrarily close to zero, the algorithms in practice become unstable and need to be turned off."," Note that while the plots for the dual-reference cancellers show the residual power as the reference interference-to-noise power ratio approaches arbitrarily close to zero, the algorithms in practice become unstable and need to be turned off."
+ This point will be reiterated where appropriate in the discussion below., This point will be reiterated where appropriate in the discussion below.
+" For the theory we have set o2,=ox.0, so that all of the output power displayed in this section is a combination of residual RFI and any reference receiver noise added during cancelling."," For the theory we have set $\sigma_S^2 = \sigma_{N_m}^2 = 0$, so that all of the output power displayed in this section is a combination of residual RFI and any reference receiver noise added during cancelling."
+ Figure 5 displays the proportion of residual power in the output signal after adaptive cancelling with MK1b and MK2b cancellers., Figure \ref{residual power v Btau 1} displays the proportion of residual power in the output signal after adaptive cancelling with MK1b and MK2b cancellers.
+ The plot shows that the added reference receiver noise averages away as the integration length is increased., The plot shows that the added reference receiver noise averages away as the integration length is increased.
+" If single canceller systems were being considered, then since any noise added during cancelling is sent to an auto-correlator, the output power would remain constant (i.e., remain at the levels shown on the left hand side of figure 5))."," If single canceller systems were being considered, then since any noise added during cancelling is sent to an auto-correlator, the output power would remain constant (i.e., remain at the levels shown on the left hand side of figure \ref{residual power v Btau 1}) )."
+" It is clear that the MKIb canceller hits a limit when it reaches the residual RFI, but that the MK2b does not."," It is clear that the MK1b canceller hits a limit when it reaches the residual RFI, but that the MK2b does not."
+" As INR, decreases the normalised output power of the MK1b canceller levels off at 1, so there is no cancelling taking place."," As $\it{INR}_r$ decreases the normalised output power of the MK1b canceller levels off at 1, so there is no cancelling taking place."
+ On the other hand the MK2 canceller continues to insert more and more reference receiver noise in an attempt to match the reference RFI to the main signal RFI., On the other hand the MK2 canceller continues to insert more and more reference receiver noise in an attempt to match the reference RFI to the main signal RFI.
+" Even though the MK2b canceller always has the larger total residual it is entirely zero-mean noise and averages out with the power,radiometric factor."," Even though the MK2b canceller always has the larger total residual power, it is entirely zero-mean noise and averages out with the radiometric factor."
+" Again the reader should note that for low INR, values the MK2 canceller can become unstable and requires an additional mechanism to turn off.", Again the reader should note that for low $\it{INR}_r$ values the MK2 canceller can become unstable and requires an additional mechanism to turn off.
+" Figure 6 shows contours of constant (normalised) output power as a function of NR, and the number of samples, AvTinz."," Figure \ref{residual power v INR} shows contours of constant (normalised) output power as a function of $\it{INR}_r$ and the number of samples, $\Delta\nu\tau_{\it{int}}$ ."
+" Figures 6aa and 6bb represent MKla and MK2a cancelling respectively, and figures 6cc and 6dd represent MK1b and MK2b cancelling respectively."," Figures \ref{residual power v INR}a a and \ref{residual power v INR}b b represent MK1a and MK2a cancelling respectively, and figures \ref{residual power v INR}c c and \ref{residual power v INR}d d represent MK1b and MK2b cancelling respectively."
+" The amount of residual power in dB is indicated by the grey scale and runs from -80 to 20 dB. It is clear from figure 6cc that a constant RFI residual remains for all INR, values after MK1b cancelling.", The amount of residual power in dB is indicated by the grey scale and runs from -80 to 20 dB. It is clear from figure \ref{residual power v INR}c c that a constant RFI residual remains for all $\it{INR}_r$ values after MK1b cancelling.
+ The dashed line indicates the approximate line where the added reference receiver noise power has averaged down to expose the non-zero residual RFI powerlevel., The dashed line indicates the approximate line where the added reference receiver noise power has averaged down to expose the non-zero residual RFI powerlevel.
+"where (he kernel is given by A(Gm.m)eo(n.imnNVGn.ain)S(m.an). with o(m.Kant? amy, Nu=z(3/Anp, and p. is the particle material density. assumed constant.","where the kernel is given by $K(m,m^\prime) = \sigma(m,m^\prime)\Delta V(m,m^\prime) S(m,m^\prime)$, with $\sigma(m,m^\prime) = K_0(m^{1/3}+m^{\prime 1/3})^2$ , $K_0 = \pi(3/4\pi\rho_s)^{2/3}$, and $\rho_s$ is the particle material density, assumed constant."
+ Note that for the special case of q=1. Eq. (," Note that for the special case of $q=1$, Eq. ("
+"18) has a shehtly different form which depends on In(my,/mg).",18) has a slightly different form which depends on $\ln{(m_L/m_0)}$.
+ After some simple algebra. E," After some simple algebra, Eq."
+s (13) and (19) can be written as which we integrate using a fourth order Iunge-Ixutta scheme.,'s (18) and (19) can be written as which we integrate using a fourth order Runge-Kutta scheme.
+ Equation (17) then gives ihe moments δι and Als : isa function o (dme. which can be directly compared with direct integration of the same coxditions using the coagulation equation (Eq.," Equation (17) then gives the moments $M_0$ and $M_2$ as a function of time, which can be directly compared with direct integration of the same conditions using the coagulation equation (Eq."
+ 1)., 1).
+" The advantage of thi5 approach (ii which a powerlaw is asstumecl at all times) is ils (ransparency: (hat is. ihe variables being sought Gr, and ο) are solved for directly."," The advantage of this approach (in which a powerlaw is assumed at all times) is its transparency; that is, the variables being sought $m_L$ and $c$ ) are solved for directly."
+ Furthermore. the change i1 (he coagulation kernel as the particle size distribution changes is included because (he kernel is updated aid expicitlv integrated into Dy and D» with every lime step.," Furthermore, the change in the coagulation kernel as the particle size distribution changes is included because the kernel is updated and explicitly integrated into $\Gamma_0$ and $\Gamma_2$ with every time step."
+ This will prove advantageous when additional effects such as sticking are included in the kernel., This will prove advantageous when additional effects such as sticking are included in the kernel.
+ In addition. source and sink terms reel not parameterized in terms of integer nomenis and can be implemented directly.," In addition, source and sink terms need not parameterized in terms of integer moments and can be implemented directly."
+ An iuilortunate disadvantage of this approach is (hat because the kernel must be integrated {11 [fact several times) over both i and m every (ime step to get Py and D». the CPU time involved is significantly longer than a fully implicit case (e.g.. 2.1: also see 23.2): however. ib remains a much faster approach (ordersof magnitude) than solving Eq. (," An unfortunate disadvantage of this approach is that because the kernel must be integrated (in fact several times) over both $m$ and $m^\prime$ every time step to get $\Gamma_0$ and $\Gamma_2$, the CPU time involved is significantly longer than a fully implicit case (e.g., 2.1; also see 2.2.2); however, it remains a much faster approach (ordersof magnitude) than solving Eq. ("
+1) directly since the cumbersome convolution has been eliminated.,1) directly since the cumbersome convolution has been eliminated.
+NGC 6791 is an extreme Galactic star cluster with an old age of ο” Gyr 2008).. a high metallicity |Fe/1I]|—4-0.30 (Doesgaardetal.2009)... and an unusual orbit that periodically brings it close to the bulge of the Milky Way (Be,"NGC 6791 is an extreme Galactic star cluster with an old age of $\sim$ 8 Gyr \citep{gru08}, a high metallicity $+$ 0.30 \citep{boe09}, and an unusual orbit that periodically brings it close to the bulge of the Milky Way \citep{bed06}."
+clinetal.2006).. was the [inst (ο provide an estimate lor the total mass of NGC 6791: 37004. down to V. =20. eonfirmed by Naluzny&Udalski(1992).," \citet{kin65} was the first to provide an estimate for the total mass of NGC 6791: $\sim$ $M_\sun$ down to $V=$ 20, confirmed by \citet{kal92}."
+. It is widely acknowledged that NGC 6191 is one of the most massive old open clusters in our Galaxy., It is widely acknowledged that NGC 6791 is one of the most massive old open clusters in our Galaxy.
+ As indicated by the discovery of several extremely blue subdwarls 1992) and (he presence of a prominent red clump. morphology of the color-magnitude diagram (CMD) for NGC 6791 is complex.," As indicated by the discovery of several extremely blue subdwarfs \citep{kal92} and the presence of a prominent red clump, morphology of the color-magnitude diagram (CMD) for NGC 6791 is complex."
+ The proposed enhanced mass loss along the red ejant branch (RGB) appears to explain the presence of hot subdwarls and the abnormally voung 2.4 Gvr white dwarl cooling age (INaliraiοἱal.2007)., The proposed enhanced mass loss along the red giant branch (RGB) appears to explain the presence of hot subdwarfs and the abnormally young 2.4 Gyr white dwarf cooling age \citep{ka07}.
+. Recently. reported that the CMD in the inner part of NGC 6791 (£2< 24) is somewhat different from that in ils outer part (2&R« 5‘).," Recently, \citet{twa11} reported that the CMD in the inner part of NGC 6791 $R<2\arcmin$ ) is somewhat different from that in its outer part $2\arcmin<R<5\arcmin$ )."
+ Specifically. there is a clichotomy near the main-sequence turnoll. which (he authors interpret as a result of protracted star formation.," Specifically, there is a dichotomy near the main-sequence turnoff, which the authors interpret as a result of protracted star formation."
+" This cluster is critically important to understanding a number of issues. such as stellar evolution and population svnthesis at high metallicity in the Milky Way and other galaxies. specifically, in ellipticals and in spirals with bulges."," This cluster is critically important to understanding a number of issues, such as stellar evolution and population synthesis at high metallicity in the Milky Way and other galaxies, specifically, in ellipticals and in spirals with bulges."
+ Due to its relative proximitv and its abundance of members. NGC 6791 provides perhaps the best opportunity [ον detailed studies of old and metal-rich populations.," Due to its relative proximity and its abundance of members, NGC 6791 provides perhaps the best opportunity for detailed studies of old and metal-rich populations."
+ Despite a rich array of extant photometric and spectroscopic studies. however. the basic knowledge of cluster membership bv. kinematic means has been limited until now to radial velocities of a few dozen bright stars (e.g..Scott and preliminary proper motions for a few thousand stars (Cudworth 2006)..," Despite a rich array of extant photometric and spectroscopic studies, however, the basic knowledge of cluster membership by kinematic means has been limited until now to radial velocities of a few dozen bright stars \citep[e.g.,][]{sco95,car06} and preliminary proper motions for a few thousand stars \citep{cud93,bed06}. ."
+ NGC 6791 is located in the field of view (FOV), NGC 6791 is located in the field of view (FOV)
+sPhe back-reaction. of. Lorentz forces. on the Ελα Mcirculation. in the convective zone of ‘a rotating --late-twpeM star‘ is of Fundamental| importancepori in theories of stellar| magnetic-field“> generation.,The back-reaction of Lorentz forces on the fluid circulation in the convective zone of a rotating late-type star is of fundamental importance in theories of stellar magnetic-field generation.
+ Over the last decade or so. a consensus has emerged. among dvnamo theorists that this back-reaction could have a significant influence on the dillerential rotation pattern in the convective zones of the Sun and active tvpe stars (see. e.g. Brandenburgetal.1991:Moss1995:Ixitchatinovetal. 1999)).," Over the last decade or so, a consensus has emerged among dynamo theorists that this back-reaction could have a significant influence on the differential rotation pattern in the convective zones of the Sun and active late-type stars (see, e.g. \pcite{brandenburg91,moss95,kitchatinov99SoPh}) )."
+ The quasi-evcelie orbital period changes observed among short-periocd binaries with magnetically-active components xovide further indirect evidence that substantial changes in differential rotation could be taking place on active binary components., The quasi-cyclic orbital period changes observed among short-period binaries with magnetically-active components provide further indirect evidence that substantial changes in differential rotation could be taking place on active binary components.
+ Work bv Applegate(1992).. Lanza.Rodono&tosner(1998). and Lanza&Rodond(1909). suggests that small. magnetically-mocdulated changes in stellar cilferential rotation will alter the gravitational quadrupole moment of he active star.," Work by \scite{applegate92}, \scite{lanza98} and \scite{lanza99} suggests that small, magnetically-modulated changes in stellar differential rotation will alter the gravitational quadrupole moment of the active star."
+ This can be sullicient to. produce orbital »eriod changes as large the fractional P7Pc1.5.10.+ observed in HU 1099 by asDonati(1999)., This can be sufficient to produce orbital period changes as large as the fractional $\delta P/P\simeq 1.5\times 10^{-4}$ observed in HR 1099 by \scite{donati99hr1099}.
+" In this paper we report observations of long-term. changes in surface cilferential rotation on the voung. rapidly rotating WO chwarl AB Doraclus. providing independent supportpper for this theory,"," In this paper we report observations of long-term changes in surface differential rotation on the young, rapidly rotating K0 dwarf AB Doradus, providing independent support for this theory."
+ We applypp? a new method for determiningo the latitudes. ancl rotation rates of⋅⊀⊀⊀ individual starspots (Collierqe'Cameron.Donati.&Semel!2002). to six. sets of archival. time-resolved. echelle spectroscopy of. AB Dor. spanning the period from 1988 December to 1996 December.," We apply a new method for determining the latitudes and rotation rates of individual starspots \cite{cameron2001diffrot} to six sets of archival time-resolved echelle spectroscopy of AB Dor, spanning the period from 1988 December to 1996 December."
+ The details of the instruments anc observing procedures used to secure the six data. sets. have been. published elsewhere. so we give only a brief summary in Table 1..," The details of the instruments and observing procedures used to secure the six data sets have been published elsewhere, so we give only a brief summary in Table \ref{tab:obser}."
+" ""he observations from the three earliest runs were re-extracted from the raw data using the Starlink ECTIIOMODP optimal extraction routines. to ensure consistency. with the ater data."," The observations from the three earliest runs were re-extracted from the raw data using the Starlink ECHOMOP optimal extraction routines, to ensure consistency with the later data."
+ The spectra from all 6 vears included. the Ha region. in which numerous weak. narrow telluric lines. of 11l;O and Os» are present.," The spectra from all 6 years included the $\alpha$ region, in which numerous weak, narrow telluric lines of $_2$ O and $_2$ are present."
+ We used. spectral subtraction (CollierCameronctal.2001). to isolate the travelling distortions produced. by starspots in the mean profile. and cast-squares deconvolution (Donatietal.1997). to stack up the residual profile information in the large number of known photospheric lines recorded in the echellograms.," We used spectral subtraction \cite{cameron2001upsand} to isolate the travelling distortions produced by starspots in the mean profile, and least-squares deconvolution \cite{donati97zdi} to stack up the residual profile information in the large number of known photospheric lines recorded in the echellograms."
+ The number IN of lines used in each vear ranges from 140 to 2000.," The number $N$ of lines used in each year ranges from 140 to 2000,"
+and the variance of the truncated Pareto pal is with The variance of the truncated Pareto is alwavs defined for every value of e> 0: conversely the variance of (he Pareto distribution can be defined only when e> 2.,and the variance of the truncated Pareto pdf is with The variance of the truncated Pareto is always defined for every value of $c >$ 0; conversely the variance of the Pareto distribution can be defined only when $c >$ 2.
+ The parameter € can be derived through the maximum likelihood estimator (AILE)., The parameter $c$ can be derived through the maximum likelihood estimator (MLE).
+ The likelihood function is defined as the probability. we would have obtained a given set of observations if given a particular set of values of the distribution parameters.e;. If we assume that the n ranclom variables ave independently aud identically distributed. then we may write the likelihood funetion as The maximum likelihood estimates for the e; are obtained by maximising the likelihood function. ο).," The likelihood function is defined as the probability we would have obtained a given set of observations if given a particular set of values of the distribution $c_i$, If we assume that the n random variables are independently and identically distributed, then we may write the likelihood function as The maximum likelihood estimates for the $c_i$ are obtained by maximising the likelihood function, L(c)."
+ Equivalently. we may find it easier to maximise /nf(Cr;). termed the So.," Equivalently, we may find it easier to maximise $ln f(x_i)$, termed the log-likelihood."
+" for a random sample 4....7, from a truncated. Pareto distribution. the likelihood function is given bv In this model we have assumed that à— min(ry....r,,) and b= max(rj..v, )."," So, for a random sample $x_1 \ldots x_n$ from a truncated Pareto distribution, the likelihood function is given by In this model we have assumed that $a$ = $x_1 \ldots x_n$ ) and $b$ = $x_1 \ldots x_n$ )."
+ Usine logarithms. we obtain the log-likelihood," Using logarithms, we obtain the log-likelihood"
+for the FIGOW image.,for the F160W image.
+ The expected noise from (he individual pixel noise described above. including the [actor of 1.8. is 3.5xLO7 which is comparable to the measured value.," The expected noise from the individual pixel noise described above, including the factor of 1.8, is $3.5 \times 10^{-3}$ which is comparable to the measured value."
+ To the degree that the aperture noises are (uly Gaussian distributed. (his indicates that [actor of 1.8 is a reasonable figure to account for the correlated noise and that the aperture noise is approximately equal to the square root of the number of pixels in the aperture limes (he individual pixel noise.," To the degree that the aperture noises are truly Gaussian distributed, this indicates that factor of 1.8 is a reasonable figure to account for the correlated noise and that the aperture noise is approximately equal to the square root of the number of pixels in the aperture times the individual pixel noise."
+" The measured aperture noise in Janskys from (he histogram is 5.8x10? which is equivalent to an AB magnitude of 29.5 for the 0.54"" diameter aperture.", The measured aperture noise in Janskys from the histogram is $5.8 \times 10^{-9}$ which is equivalent to an AB magnitude of 29.5 for the $0.54\arcsec$ diameter aperture.
+ The 5e AB magnitude is 27.7 which we will take as the appropriate value independent of source noise., The $5\sigma$ AB magnitude is 27.7 which we will take as the appropriate value independent of source noise.
+" In observations of the NIIDE-S Labbéetal.(2003) with ISAAC on the VLT found Lo aperture noises of 28.6 and 23.1 for the J and Hl bands with a 0.7"" diameter aperture which gives 5o values of 26.7 and 26.2.", In observations of the NHDF-S \citet{lab03} with ISAAC on the VLT found $1\sigma$ aperture noises of 28.6 and 28.1 for the J and H bands with a $0.7\arcsec$ diameter aperture which gives $5\sigma$ values of 26.7 and 26.2.
+ The minimum number of contiguous pixels for a real source is sel (o 7 in the extraction procedure., The minimum number of contiguous pixels for a real source is set to 7 in the extraction procedure.
+ The 5o noise for a point source detection is ihen the noise in the 7 pixel aperture which gives AD magnitudes of 30.35 and 30.15 for the F110W and FIGOW fillers respectively., The $5 \sigma$ noise for a point source detection is then the noise in the 7 pixel aperture which gives AB magnitudes of 30.35 and 30.15 for the F110W and F160W filters respectively.
+ DE in (he two image mode uses a detection image to determine the position and extent of sources., SE in the two image mode uses a detection image to determine the position and extent of sources.
+ The individual image source extraction is then performed on exactly (he same positions and regions determined [rom (he detection image., The individual image source extraction is then performed on exactly the same positions and regions determined from the detection image.
+ The SE parameters regarding source geomel(rv such as area ancl ellipticity are determined by the detection image., The SE parameters regarding source geometry such as area and ellipticity are determined by the detection image.
+ The detection image for (he treasury catalog is the simple sum of the FILOW and FIGOW science images., The detection image for the treasury catalog is the simple sum of the F110W and F160W science images.
+ Even though it is the sum of two images. the detection image has a significantly lower signal to noise than any of the ACS UDF images except for sources that are extremely red.," Even though it is the sum of two images, the detection image has a significantly lower signal to noise than any of the ACS UDF images except for sources that are extremely red."
+ Users that are interested in the NICMOS limits on faint ACS UDF sources should use ihe ACS images as the detection image to perform the source extraction on the NICMOS images., Users that are interested in the NICMOS limits on faint ACS UDF sources should use the ACS images as the detection image to perform the source extraction on the NICMOS images.
+ On the other hand very red sources may appear only in the NICMOS images., On the other hand very red sources may appear only in the NICMOS images.
+ Since we wish to provide an infrared catalog we chose to use the NICMOS images for extraction., Since we wish to provide an infrared catalog we chose to use the NICMOS images for extraction.
+ Note that by combining the two NICMOS images there is a bias against the verv recldest sources (hat might only appear in the FIGOW image., Note that by combining the two NICMOS images there is a bias against the very reddest sources that might only appear in the F160W image.
+ SE utilizes a rms image to determine the detection limit of a pixel signal when operating in the RAIS mode used in the treasury version 2.0 catalog., SE utilizes a rms image to determine the detection limit of a pixel signal when operating in the RMS mode used in the treasury version 2.0 catalog.
+ The drizzle procedure produces an observation (time weight map that measures the total integration (ime for every. pixel but the, The drizzle procedure produces an observation time weight map that measures the total integration time for every pixel but the
+observed structure.,observed structure.
+ IIST observatious of other AGNs sugeest three possible explanations for the observed clongated structure: 1) gas which is shocked aud compressed by interaction with the jet: 2) illuiiuation of gas clouds by an anisotropic nuclear radiation field G.c. an ionization conc. cf.," HST observations of other AGNs suggest three possible explanations for the observed elongated structure: 1) gas which is shocked and compressed by interaction with the jet; 2) illumination of gas clouds by an anisotropic nuclear radiation field (i.e. an ionization cone, cf."
+ Robinson1997 and references there): aud 3) a gaseous disk. such as tha of Ms? (armsetal.199 19).," \cite{robinson:97}
+ and references therein); and 3) a gaseous disk, such as that of M87 \cite{harms:94}) )."
+ The structures position angle differs bv ~207 from he jet axis (see Fie. D)., The structure's position angle differs by $\sim20\arcdeg$ from the jet axis (see Fig. \ref{fig:pacont}) ).
+ This significant misaliennieu sugeests the cussion is nof due to a current jet-cloud interaction. especially when we see other features well correlated with the vacdio/N-+vav Ίο norphology (see next section}.," This significant misalignment suggests the emission is not due to a current jet-cloud interaction, especially when we see other features well correlated with the radio/X-ray jet morphology (see next section)."
+ It las long been ivpothesized that the radio jet had a different (smaller) »osition anele in the past. creating the N-S oricutation of he outer lobes. aud then rotating to form first the ner obes and then the ciently observed X-rayradio jet.," It has long been hypothesized that the radio jet had a different (smaller) position angle in the past, creating the N-S orientation of the outer lobes, and then rotating to form first the inner lobes and then the currently observed X-ray/radio jet."
+ Iu lis regard we note that the outer edge of the N-E inner obe docs have structure aligned with the position angle of our feature: both could be gas shocked by the jet (perlaps wo-sidled} at some recent time. aud not vet cooled down.," In this regard we note that the outer edge of the N-E inner lobe does have structure aligned with the position angle of our feature; both could be gas shocked by the jet (perhaps two-sided) at some recent time, and not yet cooled down."
+ Existing radio aud N-vav measuremieuts do not provide added information ou this small spatial scale., Existing radio and X-ray measurements do not provide added information on this small spatial scale.
+ The ~20° nisalieument with the jet does allow the clongated structure to lie well within a putative ionization cone., The $\sim20\arcdeg$ misalignment with the jet does allow the elongated structure to lie well within a putative ionization cone.
+ We note that most of the other detected cussion features would also lie within such a cone if it las au opening half-augle of at least 30°. well within the liuits postulated by the unified model (c.g. Antonucci 1993)).," We note that most of the other detected emission features would also lie within such a cone if it has an opening half-angle of at least $\sim30\arcdeg$, well within the limits postulated by the unified model (e.g. \cite{antonucci:93}) )."
+ As it stands. we cannot rule out clouds deusely distributed along this position augle. embedded: in a laree opening anele radiation feld.," As it stands, we cannot rule out clouds densely distributed along this position angle, embedded in a large opening angle radiation field."
+ However. the relative thickness of the feature in the transverse direction aud the svuunetry around the nucleus sugeest that this is not the case.," However, the relative thickness of the feature in the transverse direction and the symmetry around the nucleus suggest that this is not the case."
+ We fud the simplest interpretation for the clougated enission to be a gaseous disk around the nucleus; as seen ou larger scales around other galactic nuclei," We find the simplest interpretation for the elongated emission to be a gaseous disk around the nucleus, as seen on larger scales around other galactic nuclei."
+ It could be the outer part of an accretion disk. expected around a massive black-hole at the core of this ACN.," It could be the outer part of an accretion disk, expected around a massive black-hole at the core of this AGN."
+ Ifthe structure is indeed a thin circular disk. the axial ratio suggests au inclination of 2607," If the structure is indeed a thin circular disk, the axial ratio suggests an inclination of $\simeq$."
+ Its radius of ~172 (~ 20pc) makes it sienificautly sanaller than the hundred parsec scale stellar disks observed by IIST at the ceuters of many galaxies., Its radius of $\sim 1\farcs2$ $\sim 20$ pc) makes it significantly smaller than the hundred parsec scale stellar disks observed by HST at the centers of many galaxies.
+" This eas disk could be readily ionized by the powerful ACN, seen in N- aud y-ravs."," This gas disk could be readily ionized by the powerful AGN, seen in X- and $\gamma$ -rays."
+" The flux of photous required to keep the celuitting imaterial ionized can be computed following Osterbrock(1989): after dereddening =110 mae) aud assuniue Case D recombination for sts and Ὁ, we find a value of photous s+."," The flux of photons required to keep the emitting material ionized can be computed following \cite{osterbrock:89}: after dereddening 10 mag) and assuming Case B recombination for K and $^{-3}$, we find a value of photons $^{-1}$."
+ This is a snall fraction of the total expected cussion frou the AGN., This is a small fraction of the total expected emission from the AGN.
+ Note that the nucleus has an extinction of up to TOmmae alone the line of sight. estimated from N-rav observations.," Note that the nucleus has an extinction of up to 70mag along the line of sight, estimated from X-ray observations."
+ However. while being within the ionization cone the disk is well outside the obscuring torus. aud subject to a a nmch lower extinction from the uucleus.," However, while being within the ionization cone the disk is well outside the obscuring torus, and subject to a a much lower extinction from the nucleus."
+ We thus assume —110 mag as a reasonable uppor estimate of the foreground. extinction. following Schlireier et al. (," We thus assume 10 mag as a reasonable upper estimate of the foreground extinction, following Schreier et al. ("
+1996) who found a peak extinction of —77 mag in the optical.,1996) who found a peak extinction of 7 mag in the optical.
+" The mass of the gas responsible for both exteuded aud unresolved uuclear enassion ds estimated from the standard relation. £(Pan)=Nola (Paa)Mfau where lis the protonN,, denusitv. V is the volume of the enüttiug eas, f(Paa) is the line euissivitv. nu is the proton mass. Fels(Pan) is the observed ffüux in units of ere cu2s + aud 109009610) js the reddening correction if lis different from the assumed. 10 mae."," The mass of the gas responsible for both extended and unresolved nuclear emission is estimated from the standard relation, $L(\PA) = \Np \Ne V 4\pi\cal J(\PA) = \Ne$ $\pi\cal J(\PA) {\rm M}/{\rm m_H}$: where is the proton density, V is the volume of the emitting gas, ${\cal J}(\PA)$ is the line emissivity, $\rm m_{\rm H}$ is the proton mass, $F_{-13}^{\rm obs}(\PA)$ is the observed flux in units of erg $^{-2}$ $^{-1}$ and $10^{0.059(\Av-10)}$ is the reddening correction if is different from the assumed 10 mag."
+ This modest mass estimate of M ~ depends mostly on the assunued gas deusity cm) and is extremely uncertain. but we believe that a disk of tto ds quite feasible.," This modest mass estimate of M $\sim$ depends mostly on the assumed gas density $^{-3}$ ) and is extremely uncertain, but we believe that a disk of to is quite feasible."
+ The dadisk is consistent with being perpendicular to the plaue of the dust lane. and Iwine aloug the major axis of the large elliptical ealaxy (sec Fig. 1) ," The disk is consistent with being perpendicular to the plane of the dust lane, and lying along the major axis of the large elliptical galaxy (see Fig. \ref{fig:pacont}) )."
+"Tt is now widely accepted that the laree gas/dust disk of Cou A G.c. the dust hue) was acquired in a recent mereer process. aud that differential precession has led to the observed warped. structure (οιο, Tubbs 1950))."," It is now widely accepted that the large gas/dust disk of Cen A (i.e. the dust lane) was acquired in a recent merger process, and that differential precession has led to the observed warped structure (e.g. \cite{tubbs:80}) )."
+ In the ceutral region of Ceu A. the precession time is only vvears. aud the orthogonal alieumieut of our small disk nav be consistent with unuerical studies of the evolution of eas disks in bulge systems (0.5. Quillenetal. 1992)).," In the central region of Cen A, the precession time is only years, and the orthogonal alignment of our small disk may be consistent with numerical studies of the evolution of gas disks in bulge systems (e.g. \cite{quillen:92}) )."
+ If we thus interpret the observed cussion as being from he warped outer portions of au accretion disk around the active nucleus. we couclude that even on the relatively small spatial scale of a few parsecs. the disk is dominated x the exavitational potential of the galaxy as a whole aud rot by the symmetry of the ACN and its jet.," If we thus interpret the observed emission as being from the warped outer portions of an accretion disk around the active nucleus, we conclude that even on the relatively small spatial scale of a few parsecs, the disk is dominated by the gravitational potential of the galaxy as a whole and not by the symmetry of the AGN and its jet."
+ For a non-rotating black hole (BIT). oue would expect he jet to have its direction determined by the angular uonmentun axis of the Πιο: gaseous disk. while for a I&err dack hole. it would be expected to be aligned along the spin axis of the black hole itself (Bardeen 1975).," For a non-rotating black hole (BH), one would expect the jet to have its direction determined by the angular momentum axis of the inner gaseous disk, while for a Kerr black hole, it would be expected to be aligned along the spin axis of the black hole itself \cite{bardeen:75}) )."
+ Close to a Kerr BIL ἐν~ GCun/c? ~ cou ~ pe). the disk itself will warp to become normal to the spin axis.," Close to a Kerr BH, $r\simeq$ $^2$ $\sim$ cm $\sim$ pc), the disk itself will warp to become normal to the spin axis."
+ In our Con A data. ou a few parsec scale. the jornal to the ddisk and the radio jet are misaligned by ~707. in xojection.," In our Cen A data, on a few parsec scale, the normal to the disk and the radio jet are misaligned by $\sim$, in projection."
+ We conclude that if Con A coutains a rotating jack hole. then either the gas disk must be outside the sphere of influence of the black hole. or it was formed recently enough that it has not vet become aligned with he spin axis: for a nou-rotating black hole. we fuel that the disk nist become signuificautly warped away from being jiormal to the jet inside a radius of ~2pc.," We conclude that if Cen A contains a rotating black hole, then either the gas disk must be outside the sphere of influence of the black hole, or it was formed recently enough that it has not yet become aligned with the spin axis; for a non-rotating black hole, we find that the disk must become significantly warped away from being normal to the jet inside a radius of $\sim$ 2pc."
+ Pringle(1997) has recently iodelecd selfiuduced warping of acerction disks in ACNs.," \cite{pringle:97}
+ has recently modeled self-induced warping of accretion disks in AGNs."
+ Πο finds that disks are likely to be warped at Ro> O.02pe around a ~LO* , He finds that disks are likely to be warped at $R > 0.02$ pc around a $\sim$ 
+For observations of low-excitation PN. the slit was placed in the outer edge of the O regions where the projected velocity along the hue of sight is quite low.,"For observations of low-excitation PN, the slit was placed in the outer edge of the O regions where the projected velocity along the line of sight is quite low."
+ Therefore. the narrow high-ionization lines. such as 1Ο 111 and [Ne 11]. ave dominated by thermal aud instrumental broadening.," Therefore, the narrow high-ionization lines, such as [O ] and [Ne ], are dominated by thermal and instrumental broadening."
+ The [O ni AA1363. lines are therefore unsuitable for deriving teniperature 5007.variations in velocity space.," The [O ] $\lambda\lambda4363,5007$ lines are therefore unsuitable for deriving temperature variations in velocity space."
+ We COMPALC the mwedicted ORL profiles with the observations., We compare the predicted ORL profiles with the observations.
+ Figure |0. displays a few examples for such a conrparison., Figure \ref{ic418_r} displays a few examples for such a comparison.
+ Generally. our model shows good agrecuent with observations of the hieh-photoionization lines. €11. Nou. Ot. and Nm. IC 118 has a small ORL/CEL abundance discrepancy ((O?!IIDogz/(0?! 1.3: Sharpee et al. 200801. ," Generally, our model shows good agreement with observations of the high-photoionization lines, C, N, O, and N. IC 418 has a small ORL/CEL abundance discrepancy $({\rm O}^{2+}/{\rm H}^+)_{\rm ORL}/({\rm O}^{2+}/{\rm H}^+)_{\rm CEL}=1.3$ ; Sharpee et al. \cite{sharpee04}] ],"
+and is therefore an unsuitable target fo investigate the materials in which ORLs originate., and is therefore an unsuitable target to investigate the materials in which ORLs originate.
+ Figure 10 demonstrates that our model provides a reasonable match to the O 155123. but is unable to explain the profiles of O ATTT3 and N A8223. whose widths appear to be broader that those of the predicted les.," Figure \ref{ic418_r} demonstrates that our model provides a reasonable match to the O $\lambda5513$, but is unable to explain the profiles of O $\lambda7773$ and N $\lambda8223$, whose widths appear to be broader that those of the predicted lines."
+ As pointed out bv Sharpee et al. (2001)).," As pointed out by Sharpee et al. \cite{sharpee04}) ),"
+ fluorescence excitation niv contribute substantially to the iuteusitv of these lines. aud the radiation of the lines may originate predominantly in the outer neutral regions.," fluorescence excitation may contribute substantially to the intensity of these lines, and the radiation of the lines may originate predominantly in the outer neutral regions."
+ This yaper addresses the poteutial of oenüssion line profiles to probe the physical conditions of PNe., This paper addresses the potential of emission line profiles to probe the physical conditions of PNe.
+ Our results show that if temperature or density variations are present within the nebulae. lines with differcut excitation temperatures or critical densities nieht show siguificautlv different profiles even though they originate im similar ionic species.," Our results show that if temperature or density variations are present within the nebulae, lines with different excitation temperatures or critical densities might show significantly different profiles even though they originate in similar ionic species."
+ We show that line profiles provide a new wav to investigate the CEL/ORL abuudauce-discerepaucy problem., We show that line profiles provide a new way to investigate the CEL/ORL abundance-discrepancy problem.
+ For this purpose. we require ligh-resolition aud deep spectroscopic data.," For this purpose, we require high-resolution and deep spectroscopic data."
+ In the ideal case. the line width contributed by instruuieutal broadcuing should be smaller than that bv the thermal broadening of leavy-clement lines.," In the ideal case, the line width contributed by instrumental broadening should be smaller than that by the thermal broadening of heavy-element lines."
+ This technical requirement corresponds to a spectral resolution higher than 75.000.," This technical requirement corresponds to a spectral resolution higher than $75,000$."
+ We preseut the electron temperatures and densities in the velocity space of a sample of PNe., We present the electron temperatures and densities in the velocity space of a sample of PNe.
+ No siguificaut telmperature or cdeusitv variations are found., No significant temperature or density variations are found.
+ This provides a lower lit to the temperature are clensity iuhoimosgeneities along the Ime of sight., This provides a lower limit to the temperature and density inhomogeneities along the line of sight.
+ We attempt to reproduce the observational data of two PNe. NGC 6153 and NGC 7009. acquired w Barlow ct al. (2006)).," We attempt to reproduce the observational data of two PNe, NGC 6153 and NGC 7009, acquired by Barlow et al. \cite{barlow06}) ),"
+ which show that the |O nm] A5007. the ο 111 À1363. and the ο recombination lines lave «coc.fforeut profiles.," which show that the [O ] $\lambda5007$, the [O ] $\lambda4363$, and the O recombination lines have different profiles."
+ We fud hat a pure photoionization model of chemically-homoecneous eas can explain the |O πι ADOOT /À1363 profile discrepancies. but caunot explain the [O 1u]/O profile discrepancies.," We find that a pure photoionization model of chemically-homogeneous gas can explain the [O ] $\lambda5007$ $\lambda4363$ profile discrepancies, but cannot explain the [O ]/O profile discrepancies."
+ We thus conclude that CEL and ORL may originate in kincmatically-diftercut nebular componecuts., We thus conclude that CEL and ORL may originate in kinematically-different nebular components.
+ Alternatively. the [O ui/O xofile discrepancies could be caused bv extra heating in the outer regions.," Alternatively, the [O ]/O profile discrepancies could be caused by extra heating in the outer regions."
+ Tn this work. 1D colupiutations were completed.," In this work, 1D computations were completed."
+ Line xofiles mav significautly depend on nebular geometrical structures. as proposed by Morisset Stasiisska (200623).," Line profiles may significantly depend on nebular geometrical structures, as proposed by Morisset Stasińsska \cite{morisset06a}) )."
+ Tlowever. for lines from tle same ionic species. nebular ecometrical structures provide an effect ou the profiles of a similar degree. and thus hardly affect our results.," However, for lines from the same ionic species, nebular geometrical structures provide an effect on the profiles of a similar degree, and thus hardly affect our results."
+ We coustruct line profiles of the approximatively-spherical PN. IC 11δ.," We construct line profiles of the approximatively-spherical PN, IC 418."
+ Our model can explain most of he observed Hue profiles., Our model can explain most of the observed line profiles.
+ A velocity feld that sharply lnereases outwards is revealed., A velocity field that sharply increases outwards is revealed.
+ We fud that the |S ui] and [O n] deusity diagnostic lines have very different xofiles. which imply that large deusitv variations are xesent within the PN.," We find that the [S ] and [O ] density diagnostic lines have very different profiles, which imply that large density variations are present within the PN."
+ Our model shows that in the low-velocity reeious of the |O | AG300 line. the predicted fux is ower than the observed one. which indicates the existence of neutral chuups within the ionized regions.," Our model shows that in the low-velocity regions of the [O ] $\lambda6300$ line, the predicted flux is lower than the observed one, which indicates the existence of neutral clumps within the ionized regions."
+ Concerally. he profiles of CELs and ORLs are in good agreement.," Generally, the profiles of CELs and ORLs are in good agreement."
+" However, IC. 118 is a τοις PN which has a small CEL/ORL abuudaucediscrepaucy. and thus is uusuitable or studving the abundance problem."," However, IC 418 is a young PN which has a small CEL/ORL abundance-discrepancy, and thus is unsuitable for studying the abundance problem."
+ For further studies. Hel-quality spectroscopic observationsof PNe with large CEL/ORL abunedance-discrepancics should be invaluable.," For further studies, high-quality spectroscopic observationsof PNe with large CEL/ORL abundance-discrepancies should be invaluable."
+The analvsis of presented in (his paper is based on the radio observations conducted with the GMRT and VLA. as well as on the archival VLA data from the NRAO 1400 NMIIz FIRST and NVSS survevs (Beckeretal.1995:Condon1993).,"The analysis of presented in this paper is based on the radio observations conducted with the GMRT and VLA, as well as on the archival VLA data from the NRAO $1400$ MHz FIRST and NVSS surveys \citep{bec95,con98}."
+. The observing log for the new GAIRT and. VLA observations is given in 22. including (he name of the telescope ancl (he array configuration. the lrequency of observations ancl the primary beamwidth. as well as the dates of the exposures.," The observing log for the new GMRT and VLA observations is given in 2, including the name of the telescope and the array configuration, the frequency of observations and the primary beamwidth, as well as the dates of the exposures."
+ The 4860.1 MMIIz VLA image ol the entire source is presented in H1., The $4860.1$ MHz VLA image of the entire source is presented in 1.
+ The enlarged GAIRT and VLA images of the northern and southern lobes at four different observing [frequencies are shown in 22., The enlarged GMRT and VLA images of the northern and southern lobes at four different observing frequencies are shown in 2.
+ The images at. 1400 MMIIz are reproduced from the FIRST survey. and the instrumental characteristics for all the radio maps included in this paper are collected in 33.," The images at $1400$ MHz are reproduced from the FIRST survey, and the instrumental characteristics for all the radio maps included in this paper are collected in 3."
+ The obtained radio fluxes of and the resulting luminosities are listed in 44., The obtained radio fluxes of and the resulting luminosities are listed in 4.
+ The low-lrequency GMBRT observations al 328.8 MMITz and 617.3 AIMIz were made in (he standard manner. with each observation of the target source interspersed wilh observations of calibrator sources.," The low-frequency GMRT observations at $328.8$ MHz and $617.3$ MHz were made in the standard manner, with each observation of the target source interspersed with observations of calibrator sources."
+ The phase calibrators PINS 114162067 2298) and PINS 114424101 were observed al 328.8 MMITz ane 617.3 MAUIs. respectively. after each of several 20 exposures of the target. centered on (he core position.," The phase calibrators PKS $+$ 067 298) and PKS $+$ 101 were observed at $328.8$ MHz and $617.3$ MHz, respectively, after each of several 20 minute-long exposures of the target centered on the core position."
+ 2286 was used as the flux density and bandpass calibrator based on the scale of Baersetal.(1977)., 286 was used as the flux density and bandpass calibrator based on the scale of \citet{baa77}.
+. The total observing limes on the (arget sources were about 200 min (at 328.8 MMIIZ) and 180 min (at 617.3 MMITZ)., The total observing times on the target sources were about 200 min (at $328.8$ MHz) and 180 min (at $617.3$ MHz).
+ The data were edited and reduced using the NRAO package. and then self-calibrated to produce the best possible images.," The data were edited and reduced using the NRAO package, and then self-calibrated to produce the best possible images."
+ The high-frequency 4860.1 MMIIZ observations conducted with the VLA in its D configuration were made in two runs aimed al a different lobe of each (see 22)., The high-frequency $4860.1$ MHz observations conducted with the VLA in its D configuration were made in two runs aimed at a different lobe of each (see 2).
+ The interferometric phases were calibrated every 20 min with the phase calibrator PINS —018., The interferometric phases were calibrated every 20 min with the phase calibrator PKS $-$ 018.
+ The source 2286 was used as the primary flix density calibrator., The source 286 was used as the primary flux density calibrator.
+" 41 and 93 minute-long exposures were conducted for the fields centered on the northern and the southern lobes. respectively,"," 41 and 93 minute-long exposures were conducted for the fields centered on the northern and the southern lobes, respectively."
+ This allowed us to reach an ims noise values of about 0.048 and + for the two exposures respectivelv., This allowed us to reach an rms noise values of about 0.048 and $^{-1}$ for the two exposures respectively.
+ The data were edited. reduced and self-calibrated with the NRAO package.," The data were edited, reduced and self-calibrated with the NRAO package."
+ Images produced from the resulting datasets were corrected [or the primary beam pattern., Images produced from the resulting datasets were corrected for the primary beam pattern.
+ The three DDRGs selected for this project are J14532-3308. —3216. and JOO41+3224.," The three DDRGs selected for this project are J1453+3308, $-$ 3216, and J0041+3224."
+ The archival data lor (hese used in our modeling are summarized in 55., The archival data for these used in our modeling are summarized in 5.
+iis a highly unusual X-ray binary. whose nature has been a puzzle for many vears.,"is a highly unusual X-ray binary, whose nature has been a puzzle for many years."
+ Since the carly LOTOs. Cir N-1 has shown erratic X-ray properties: its light curve cillered dramatically cach time it was observed.," Since the early 1970s, Cir X-1 has shown erratic X-ray properties: its light curve differed dramatically each time it was observed."
+ Nevertheless. there isa periodic modulation of 16.6 d (Ixaluzienskietal.1976).. which is believed to be the orbital period of the binary.," Nevertheless, there is a periodic modulation of 16.6 d \cite{khbs76}, which is believed to be the orbital period of the binary."
+ The compact star in Cir N-1 is à neutron star. inferred from the Type E X-ray bursts observed in à brief episode Clennant. Fabian Shaler 1986)).," The compact star in Cir X-1 is a neutron star, inferred from the Type I X-ray bursts observed in a brief episode (Tennant, Fabian Shafer \nocite{tfs86b}) )."
+ Further support is provided by the Z-source behaviour seen recently in X-ray timing data from theXTE satellite (Shirev. Bradt Levine 1999)).," Further support is provided by the Z-source behaviour seen recently in X-ray timing data from the satellite (Shirey, Bradt Levine \nocite{sbl99}) )."
+ The radio counterpart of Cir N-1 is located 25 [from the centre of the supernova remnant (321.9 0.3. and. is apparently connected to the remnant by a radio nebula (Llavnesetal.1986).," The radio counterpart of Cir X-1 is located $25'$ from the centre of the supernova remnant $-$ 0.3, and is apparently connected to the remnant by a radio nebula \cite{hkl+86}."
+.. Lt shows Uares at the same period as the N-rav modulation (Llavnesetal.1978)., It shows flares at the same period as the X-ray modulation \cite{hjm+78}.
+. Cir N-1 has two areminute-scale radio jets (Stewartctal.1993)... and an arcsecond-scale asymmetric jet suggests the presence of relativistic outflow: from the source (Fenderetal.1998).," Cir X-1 has two arcminute-scale radio jets \cite{schn93}, and an arcsecond-scale asymmetric jet suggests the presence of relativistic outflow from the source \cite{fst+98}."
+. The close association ofCir X-1 with the supernova remnant suggests that the system mav be a voung (<107 v old) runaway system from a supernova explosion (Stewartctal. 1993)., The close association of Cir X-1 with the supernova remnant suggests that the system may be a young $< 10^5$ y old) runaway system from a supernova explosion \cite{schn93}.
+ The optical counterpart to Cir N-1 was identified. as a hiehly-redeenec star with strong Ho. eemission. (Whelanetal. L977)., The optical counterpart to Cir X-1 was identified as a highly-reddened star with strong $\alpha$ emission \cite{wmw+77}.
+. Ehis object was later shown to consist of three E . n ⊳∖↿⋜⊔⋅⊳∖∖∖⊽∐↓⊔⊔⋜↧↓⋅⋯⇂⋯⊳∖∪⇂↓⋅⋅↱≻⋅↱≻⊳↿↓⊔⋅⊳∖⋯∐↓↥⋖⊾↓⋅⊔⊔↓∪⊳∖⇂∪⇂∖∖⋎⋯∼↓↥↓⊳∖ ⋅⋠⋠ the true counterpart (Aloneti1992:Duncanetal.1993).," This object was later shown to consist of three stars within a radius of 5, the southernmost of which is the true counterpart \cite{mon92,dsh93}."
+. In 1997. we obtained spectra of Cir N-1 near apastron'.. using the Anglo-Australian Telescope (Johnston. Fender Wu 1999. hereafter Paper L)..," In 1997, we obtained spectra of Cir X-1 near , using the Anglo-Australian Telescope (Johnston, Fender Wu 1999, hereafter Paper \nocite{jfw99}."
+ We detected an asvnunetric llo line which can be decomposed into two components. a narrow one and abroad component which is blue-shifted with respect to the narrow component.," We detected an asymmetric $\alpha$ line which can be decomposed into two components, a narrow one and abroad component which is blue-shifted with respect to the narrow component."
+ Comparison with, Comparison with
+disk.,disk.
+ There is a considerable literature on the lanuching of clectromaguetically driven jets/winds from the surface of an accretion disk. uncer the assumption that the disk is threaded by a suitably coufigured poloidal fick DDBlaudford Pavue 1982: EKóunigl 1989: Pelletier Pucitz 1992). but little discussion of how the fick configuration was set up in the first place.," There is a considerable literature on the launching of electromagnetically driven jets/winds from the surface of an accretion disk, under the assumption that the disk is threaded by a suitably configured poloidal field Blandford Payne 1982; Könnigl 1989; Pelletier Pudritz 1992), but little discussion of how the field configuration was set up in the first place."
+ The hope appears to be that a ecneral external poloidal fick will be advected isvards bv the accretion flow in the disk., The hope appears to be that a general external poloidal field will be advected inwards by the accretion flow in the disk.
+ Tlowever. it has been demoustrated that in a standard viscously driven accretion disk. uuless the ratio of magnetic diffusivity to viscosity (the inverse maguoetic Praudtl uuniber) in the disk is very small (comparable to II/R). such iid. advection of field does not take place (Lubow. Papaloizou. Pringle 199la: Reves-Ruiz Stepinski 1996: Tlevvacrts. Priest. Bardou 1996).," However, it has been demonstrated that in a standard viscously driven accretion disk, unless the ratio of magnetic diffusivity to viscosity (the inverse magnetic Prandtl number) in the disk is very small (comparable to $H/R$ ), such inward advection of field does not take place (Lubow, Papaloizou, Pringle 1994a; Reyes-Ruiz Stepinski 1996; Heyvaerts, Priest, Bardou 1996)."
+ It is important to note. however. that iu a disk iu which the angular momentum transport is mainly duc to the effects of sclfsustaining lyvdromaguctic turbulence (Balbus ILwlev 1991: Tlawley. Camuie. Balbus 1995. 1996: Stone. ILhuwvlev. Cammue Balbus 1996: Draudeuburs. Nordluud. Steiu. Torkelsson 1995). the magnetic Praudtl uuuber is likely to be of order uuitv (Parker 1971: Pouquet. Frisch Léoorat 1976: Zeldovich. Ruzimaikin. Sokoloff 1983: Cauuto Battaglia 1985).," It is important to note, however, that in a disk in which the angular momentum transport is mainly due to the effects of self-sustaining hydromagnetic turbulence (Balbus Hawley 1991; Hawley, Gammie, Balbus 1995, 1996; Stone, Hawley, Gammie Balbus 1996; Brandenburg, Nordlund, Stein, Torkelsson 1995), the magnetic Prandtl number is likely to be of order unity (Parker 1971; Pouquet, Frisch Léoorat 1976; Zel'dovich, Ruzmaikin, Sokoloff 1983; Canuto Battaglia 1988)."
+ Thus. if the disk surrounding the black hole is at any radius a standard accretion disk iu which the douünaut mode of augular monienutuui loss is by viscous transport within the disk. poloidal magnetic flux caunot be simply advected dmwards frou mfiity.," Thus, if the disk surrounding the black hole is at any radius a standard accretion disk in which the dominant mode of angular momentum loss is by viscous transport within the disk, poloidal magnetic flux cannot be simply advected inwards from infinity."
+ Under the assuniption that the inflow of material to the dack hole is from a standard Shakira Suuvaev (1973) accretion disk. we now cousider the likely streneth of the xloidal field in the inner regions.," Under the assumption that the inflow of material to the black hole is from a standard Shakura Sunyaev (1973) accretion disk, we now consider the likely strength of the poloidal field in the inner regions."
+ Chosh Abramowicz (1997). while noting that the usual assuiptiou is that he feld streneth is eiven by equating the imagnetic xesure fas to the maxima pressure du the disk.  MModerski Sisora 1996b). draw attention o the fact that the relevant field is the one produced X dvauno processes in the disk. and that current ποΊσα simulations indicate that the limiting value of lis field Paus is such the the maguetic pressure Z4. is only a small fraction of the eas pressure in the disk.," Ghosh Abramowicz (1997), while noting that the usual assumption is that the field strength is given by equating the magnetic pressure $P_{\rm mag}$ to the maximum pressure in the disk, $P_{\rm d}$ Moderski Sikora 1996b), draw attention to the fact that the relevant field is the one produced by dynamo processes in the disk, and that current numerical simulations indicate that the limiting value of this field $B_{\rm dynamo}$ is such the the magnetic pressure $P_{\rm mag}$ is only a small fraction of the gas pressure in the disk."
+ Using this estimate for the field streneth. coupled with he aremment that the field threading the hole caunot ercatly exceed this. enables them to couclude that the ver of the DBlaudford-Zuajek process las been seriously overestimated.," Using this estimate for the field strength, coupled with the argument that the field threading the hole cannot greatly exceed this, enables them to conclude that the power of the Blandford-Znajek process has been seriously overestimated."
+ However. it should also be noted that a naenetic dynamo moechanisum m au accretion disk will teud o produce maenetic field primarily on scales of the order of the disk thickness. £7. whereas in the computations iu he Chosh Abramowiez paper they make the assumption (Guspired by the fact that in the boxes used im magnetic dvuamo simulations f7/R~1: see their Figure 1) that he fields so produced have typical leugth-scales of order Ra.," However, it should also be noted that a magnetic dynamo mechanism in an accretion disk will tend to produce magnetic field primarily on scales of the order of the disk thickness, $H$, whereas in the computations in the Ghosh Abramowicz paper they make the assumption (inspired by the fact that in the boxes used in magnetic dynamo simulations $H/R\sim1$; see their Figure 1) that the fields so produced have typical length-scales of order $R_{\rm d}$."
+ Tudeed. in the solution they present for computing the field threading the hole. they male the asstuption that he field threading the disk has a configuration such that BriBe=7/3.," Indeed, in the solution they present for computing the field threading the hole, they make the assumption that the field threading the disk has a configuration such that $B_R/B_z=7/3$."
+ However. if. as they assume. this is indeed an accretion disk with maeuetic Prandtl uuuboer of order unity. then BrD. should uot exceed a value of order ΓΗ (Lubow et al.," However, if, as they assume, this is indeed an accretion disk with magnetic Prandtl number of order unity, then $B_R/B_z$ should not exceed a value of order $H/R$ (Lubow et al."
+ 1991a)., 1994a).
+ We euipliasize that it is portant to iae the distinction between the mean (large-scale) feld id the (iall-scale) fields., We emphasize that it is important to make the distinction between the mean (large-scale) field and the (small-scale) fields.
+ Tn realitv. it seems likely that a dynanio mechauisui iu an acerction disk might be able to give rise to more elobal fields with typical leneth-scales larger than 141. mt the processes by which it can do so are as vot arecly unexplored in uunuerical simulations (sce. however. Annitage 1998).," In reality, it seems likely that a dynamo mechanism in an accretion disk might be able to give rise to more global fields with typical length-scales larger than $H$, but the processes by which it can do so are as yet largely unexplored in numerical simulations (see, however, Armitage 1998)."
+ The reason for this is that the simulations je so far been Blnuited to a small clement of the disk. and are incapable of considering the disk as a global entity.," The reason for this is that the simulations have so far been limited to a small element of the disk, and are incapable of considering the disk as a global entity."
+ Tout Pringle (1996) have sugeested that outside he main body of the disk. the internal dynamo might © able to produce magnetic loops mainly on scales of order Z7. which could then iuteract by differeutial rotation and reconnection to produce an inverse cascade to larecr chethescales (see also Romanova ct al.," Tout Pringle (1996) have suggested that outside the main body of the disk, the internal dynamo might be able to produce magnetic loops mainly on scales of order $H$, which could then interact by differential rotation and reconnection to produce an inverse cascade to larger length-scales (see also Romanova et al."
+ 1998)., 1998).
+ For the muaticular idealized case which they cousicered. they found hat BA)XAÀ. where B(A) is the flux density at scales of A or ereater.," For the particular idealized case which they considered, they found that $B(\lambda)\propto\lambda^{-1}$, where $B(\lambda)$ is the flux density at scales of $\lambda$ or greater."
+ A similay scaling melt be expected from any dyuauuo process operating through an inverse cascade., A similar scaling might be expected from any dynamo process operating through an inverse cascade.
+ Iu this case. we would expect the size of the large-scale field threadiug the disk. Bia. to be given approsimately by Noting that the energv dissipated im the disk by the dynamo mechanisin is given approximately by where the first term in parentheses represents the magnetic stress in the disk. we see that the ratio of the maximal clectromaguctic flux from the disk (eq. |," In this case, we would expect the size of the large-scale field threading the disk, $B_{\rm pd}$, to be given approximately by Noting that the energy dissipated in the disk by the dynamo mechanism is given approximately by where the first term in parentheses represents the magnetic stress in the disk, we see that the ratio of the maximal electromagnetic flux from the disk (eq. ["
+6|) to the energy generated in the disk is approximately which. using the above estimate. gives This estimate is consistent with the observations of jets and winds produced from ai variety of objects PPrinele 1993: Livio 1997).,"6]) to the energy generated in the disk is approximately which, using the above estimate, gives This estimate is consistent with the observations of jets and winds produced from a variety of objects Pringle 1993; Livio 1997)."
+ Equation (9) therefore medieates that in the case of a standard disk even the calculation of Chosh Abramowicz (1997) overcstimates the power of the Dlaudford-Zuajek process., Equation (9) therefore indicates that in the case of a standard disk even the calculation of Ghosh Abramowicz (1997) overestimates the power of the Blandford-Znajek process.
+ We have seeu frou the above that iu order for the Dlaudford-Zuajek mechanisi to dominate the power output it is necessary for the magnetic flux on open field lines threading the black hole to ereatly exceed the magnetic flux on open feld lines threading the inner regions of the disk., We have seen from the above that in order for the Blandford-Znajek mechanism to dominate the power output it is necessary for the magnetic flux on open field lines threading the black hole to greatly exceed the magnetic flux on open field lines threading the inner regions of the disk.
+ We have also seen that. if the flow," We have also seen that, if the flow"
+~0.2 aremin off from the optical axis of the X-ray telescope.,$\sim 0.2$ arcmin off from the optical axis of the X-ray telescope.
+ In the December 1999 data. it appears to be ~0.3 aremin off-axis which. in both cases. is small enough to have a negligible effect for the distortion of the PSF relative to an on-axis observation.," In the December 1999 data, it appears to be $\sim 0.3$ arcmin off-axis which, in both cases, is small enough to have a negligible effect for the distortion of the PSF relative to an on-axis observation."
+ In order to increase the precision required for accurate astrometric measurements. systematic uncertainties need to be corrected.," In order to increase the precision required for accurate astrometric measurements, systematic uncertainties need to be corrected."
+ Apart from the aspect offset correction we also considered the errors introduced by determining the event positions., Apart from the aspect offset correction we also considered the errors introduced by determining the event positions.
+ The later included corrections of the tap-ringing distortion in the position reconstruction and the correction of errors introduced by determining the centroid of the charge cloud., The later included corrections of the tap-ringing distortion in the position reconstruction and the correction of errors introduced by determining the centroid of the charge cloud.
+ Instead of using the fully processed pipeline products. we started our data reduction with level-1 files to be able to correct for these systematic effects.," Instead of using the fully processed pipeline products, we started our data reduction with level-1 files to be able to correct for these systematic effects."
+ All the data processings were performed with CIAO 3.2.1., All the data processings were performed with CIAO 3.2.1.
+ Details of the applied corrections are described in the following., Details of the applied corrections are described in the following.
+ Instabilities in the HRC electronics can lead to a tap-ringing distortion in the position reconstruction of events., Instabilities in the HRC electronics can lead to a tap-ringing distortion in the position reconstruction of events.
+ Correction has been applied to minimize this distortions in standard HRC level-1 processing which required to know the values of the amplitude scale factor (AMP_SSF)., Correction has been applied to minimize this distortions in standard HRC level-1 processing which required to know the values of the amplitude scale factor SF).
+ Such values are found in the HRC telemetry and are different for each event., Such values are found in the HRC telemetry and are different for each event.
+ Unfortunately. they are often telemetered incorrectly.," Unfortunately, they are often telemetered incorrectly."
+ In order to fix this anomaly. we followed a thread in CIAO to deduce the correct values of AMP_SSF in the level-] event file from other HRC event data and applied these corrected values to minimize the distortion.," In order to fix this anomaly, we followed a thread in CIAO to deduce the correct values of SF in the level-1 event file from other HRC event data and applied these corrected values to minimize the distortion."
+ The de-gap correction was applied to the event files in order to compensate the systematic errors introduced in the event positions by the algorithm which determines the centroid of the charge cloud exiting the HRC rear micro-channel plate., The de-gap correction was applied to the event files in order to compensate the systematic errors introduced in the event positions by the algorithm which determines the centroid of the charge cloud exiting the HRC rear micro-channel plate.
+ After correcting these systematic errors we generated the level-2 event lists files which were used thoroughly for the remaining analysis., After correcting these systematic errors we generated the level-2 event lists files which were used thoroughly for the remaining analysis.
+ We created HRC-I images of for both epochs with a binning factor of 1 so that each pixel has a side length of 0.13187 aresec., We created HRC-I images of for both epochs with a binning factor of 1 so that each pixel has a side length of 0.13187 arcsec.
+ To be able to correct for pointing uncertainties by using X- counterparts of stars which have their position known with high accuracy. we applied a wavelet source detection algorithm to the HRC-I images.," To be able to correct for pointing uncertainties by using X-ray counterparts of stars which have their position known with high accuracy, we applied a wavelet source detection algorithm to the HRC-I images."
+ Two X-ray point sources with a count rate of about and relative to that of were detected serendipitously at about ~2.5 aremin and ~5.5 aremin distance fromJ0822—4300., Two X-ray point sources with a count rate of about and relative to that of were detected serendipitously at about $\sim2.5$ arcmin and $\sim5.5$ arcmin distance from.
+. Figure | shows a 9.57 aremin field surrounding as seen with the HRC-I in April 2005., Figure 1 shows a $9.5\times 7$ arcmin field surrounding as seen with the HRC-I in April 2005.
+ Both serendipitous sources. denoted as A and B. are indicated in this figure.," Both serendipitous sources, denoted as A and B, are indicated in this figure."
+ Their X-ray properties are summarized in Table I., Their X-ray properties are summarized in Table 1.
+ In order to determine their X-ray positions with higher accuracy than possible with the wavelet analysis. we fitted a 2-D Gaussian model with the modeled PSF as a convolution kernel to both sources A and B. These fits require some information on the source energy spectrum which is not available from the HRC-I data.," In order to determine their X-ray positions with higher accuracy than possible with the wavelet analysis, we fitted a 2-D Gaussian model with the modeled PSF as a convolution kernel to both sources A and B. These fits require some information on the source energy spectrum which is not available from the HRC-I data."
+ We therefore checked the archival XMM-Newton data for both sources and found from an spectro-imaging analysis of MOSI/2 data (cf., We therefore checked the archival XMM-Newton data for both sources and found from an spectro-imaging analysis of MOS1/2 data (cf.
+ Figure | in Hui Becker 2006) that the hardness-ratios of source A and B are comparable to that of which has its energy peak at ~1.5 keV. We therefore extract the Chandra HRC-I PSF model images at 1.5 keV with the corresponding off-axis angles from the CALDB standard library files (FI) by interpolating within the energy and off-axis grids using the CIAO tool MKPSF., Figure 1 in Hui Becker 2006) that the hardness-ratios of source A and B are comparable to that of which has its energy peak at $\sim 1.5$ keV. We therefore extract the Chandra HRC-I PSF model images at 1.5 keV with the corresponding off-axis angles from the CALDB standard library files (F1) by interpolating within the energy and off-axis grids using the CIAO tool MKPSF.
+ The exposure maps were also generated at this energy by using MKEXPMAP., The exposure maps were also generated at this energy by using MKEXPMAP.
+ The size of the |—«c error circles of source A obtained by this method are 0.16 aresee and 0.07 aresee for the 1999 and 2005 observations. respectively.," The size of the $1-\sigma$ error circles of source A obtained by this method are 0.16 arcsec and 0.07 arcsec for the 1999 and 2005 observations, respectively."
+ The relatively large error i1 the December 1999 observation ts due to its shorter integrattor time and thus smaller photon statistics., The relatively large error in the December 1999 observation is due to its shorter integration time and thus smaller photon statistics.
+ The large off-axis angle (~5.5 aremin) of source B causes amarked blurring of the PSF encircled energy radius =4 aresec)., The large off-axis angle $\sim 5.5$ arcmin) of source B causes a marked blurring of the PSF encircled energy radius $\simeq 4$ arcsec).
+ Such distortion makes the source appear to be very dispersed., Such distortion makes the source appear to be very dispersed.
+ Given the patchy and uneven supernova remnant emission this source is surrounded by and the limited photor statistics we dic not succeed in obtaining its coordinates more accurate than possible with the wavelet algorithm (which ts ~0.2 aresec)., Given the patchy and uneven supernova remnant emission this source is surrounded by and the limited photon statistics we did not succeed in obtaining its coordinates more accurate than possible with the wavelet algorithm (which is $\sim 0.2$ arcsec).
+ This leaves source A as the only reference star to perform astrometric correction., This leaves source A as the only reference star to perform astrometric correction.
+ Correlating the source. position of source A with the Two Micron All Sky Survey (2MASS) catalog (Skrutskie et al., Correlating the source position of source A with the Two Micron All Sky Survey (2MASS) catalog (Skrutskie et al.
+ 2006) identified the star with the source designation 08214628—4302037 as a possible optical counterpart., 2006) identified the star with the source designation $-$ 4302037 as a possible optical counterpart.
+ Since the next nearest optical source is located about 5 aresec away from the X-ray position of source A. we adopt the 2MASS source 08214628—4302037 as its optical counterpart.," Since the next nearest optical source is located about 5 arcsec away from the X-ray position of source A, we adopt the 2MASS source $-$ 4302037 as its optical counterpart."
+ The visual magnitudes of the object are J=12.16]+ 0.027. H=11.675£0.023 and K=11.558+0.024.," The visual magnitudes of the object are $J=12.161\pm 0.027$ , $H=11.675\pm 0.023$ and $K=11.558\pm 0.024$."
+" Since its spectral type is not known with certainty. we adopted a typical X-ray-to-optical flux ratio of log(F,/F,,,)~—2.46 for stars from Krautter et al. ("," Since its spectral type is not known with certainty, we adopted a typical X-ray-to-optical flux ratio of $(F_x/F_{opt})\simeq -2.46$ for stars from Krautter et al. ("
+1999).,1999).
+ Assuming a Raymond-Smith thermal plasma model with KT=0.15 keV. jj=4x107! em (Hui Becker 2006) and solar abundances for the star's spectrum we estimated with the aid of PIMMS (version 3.8a2) its HRC-I count rate to be ~3x107 ets/s. This is in good agreement with the observed count rate of source A (cf.," Assuming a Raymond-Smith thermal plasma model with $kT=0.15$ keV, $n_H=4\times10^{21}$ $^{-2}$ (Hui Becker 2006) and solar abundances for the star's spectrum we estimated with the aid of PIMMS (version 3.8a2) its HRC-I count rate to be $\sim 3\times10^{-3}$ cts/s. This is in good agreement with the observed count rate of source A (cf."
+ Table 1)., Table 1).
+ In order to use the optical identification of the serendipitous X-ray source A as a reference source for the offset correction. we have to check whether itself shows a proper motion.," In order to use the optical identification of the serendipitous X-ray source A as a reference source for the offset correction, we have to check whether itself shows a proper motion."
+ To investigate this. we correlated its 2MASS position with the UCAC? catalog (Zacharias et al.," To investigate this, we correlated its 2MASS position with the UCAC2 catalog (Zacharias et al."
+ 2003)., 2003).
+ We unambiguously found a source with the UCAC?2 designation 13302738 as a counterpart of the X-ray source A. According to this catalog. this source has a proper motion of pra=—16.0+5.2 mas/yr. Hase=—1.7+5.2 mas/yr.," We unambiguously found a source with the UCAC2 designation 13302738 as a counterpart of the X-ray source A. According to this catalog, this source has a proper motion of $\mu_{\rm RA}=-16.0\pm5.2$ mas/yr, $\mu_{\rm dec}=-1.7\pm5.2$ mas/yr."
+ We attempted to make an independent estimate by analysing the images from the first and the second Digitised SkySurveys!., We attempted to make an independent estimate by analysing the images from the first and the second Digitised Sky.
+. From the observation dates specified for the DSS-] and DSS-2 images. we found that the epoches of these two images are separated by 5134 days.," From the observation dates specified for the DSS-1 and DSS-2 images, we found that the epoches of these two images are separated by 5134 days."
+We took four bright stars within | aremin neighbourhood of the X-ray source A,We took four bright stars within 1 arcmin neighbourhood of the X-ray source A
+The ability to conduct wide-area surveys of molecular clouds has shown that most stars form in clusters containing some hundreds to thousands of stars (?;; ?;; ?)).,The ability to conduct wide-area surveys of molecular clouds has shown that most stars form in clusters containing some hundreds to thousands of stars \citealt{Ladaetal1991}; ; \citealt{ClaBonHil2000}; \citealt{LadLad2003}) ).
+" At the same time, mid-infrared surveys such as Spitzer have shown that significant numbers of stars form in a more distributed mode (???7).."," At the same time, mid-infrared surveys such as Spitzer have shown that significant numbers of stars form in a more distributed mode \citep{AllenMegeathetal2007, Gutermuthetal2008, Gutermuthetal2009, Evansetal2009}."
+" The reason why such different modes of star formation exist, and in the same cloud OOrion A) is unclear."," The reason why such different modes of star formation exist, and in the same cloud Orion A) is unclear."
+" There has also been considerable interest as to why star formation appears to be inefficient (?),, with only a few percent of a molecular cloud's mass being turned into stars per free-fall time."," There has also been considerable interest as to why star formation appears to be inefficient \citep{Evansetal2009}, with only a few percent of a molecular cloud's mass being turned into stars per free-fall time."
+" This could imply that star formation is a slow process (?) or that it is an inherently inefficient process, but proceeds on the local dynamical timescale."," This could imply that star formation is a slow process \citep{KruTan2007} or that it is an inherently inefficient process, but proceeds on the local dynamical timescale."
+ In the latter case the efficiency must increase on small scales where bound clusters are formed., In the latter case the efficiency must increase on small scales where bound clusters are formed.
+" For example, the Orion Nebula Cluster has a median age of z10° years and a dynamical time of z:3x10? years (?).."," For example, the Orion Nebula Cluster has a median age of $\approx 10^6$ years and a dynamical time of $\approx 3 \times 10^5$ years \citep{HilHar1998}."
+ Given an overall star formation efficiency of ~50% this implies a star formation per free-fall time of1596., Given an overall star formation efficiency of $\approx 50$ this implies a star formation per free-fall time of.
+". Considering that the initial cluster cloud is likely to have been at least a factor of 2 larger (?),, this implies an efficiency of star formation per free-fall time of close to 50%."," Considering that the initial pre-cluster cloud is likely to have been at least a factor of 2 larger \citep{BonBatVin2003}, this implies an efficiency of star formation per free-fall time of close to 50."
+". To date, numerical simulations have generally chosen spherically symmetric or period boxes initial conditions of gravitationally bound clouds which collapse and fragment to form stellar clusters (???).."," To date, numerical simulations have generally chosen spherically symmetric or period boxes initial conditions of gravitationally bound clouds which collapse and fragment to form stellar clusters \citep*{KleBurBat1998,BatBonBro2003,Bate2009a}."
+ Cluster formation proceeds through hierarchical fragmentation and production of a somewhat distributed population which undergoes a hierachical merger process from small subclusters to one final cluster containing most of the stars (??;; 7)).," Cluster formation proceeds through hierarchical fragmentation and production of a somewhat distributed population which undergoes a hierachical merger process from small subclusters to one final cluster containing most of the stars \citealt*{BonBatVin2003, Bate2009a}; \citealt{Federrathetal2010}) )."
+" One simple possibility is that if star formation occurs in regions of molecular clouds that are globally unbound, then there is no reason for the stars that form from the fragmenting population to fall together to form the large stellar cluster."," One simple possibility is that if star formation occurs in regions of molecular clouds that are globally unbound, then there is no reason for the stars that form from the fragmenting population to fall together to form the large stellar cluster."
+" Recent work evaluating the boundness of molecular clouds show that their masses are typically five times smaller thanthat to be virialised, implying that much of the present day"," Recent work evaluating the boundness of molecular clouds show that their masses are typically five times smaller thanthat to be virialised, implying that much of the present day"
+ CH;OH ts formed from H:CO after CH;CHO dissociation.,$_3$ OH is formed from $_2$ CO after $_3$ CHO dissociation.
+" However. the CH:OH yield ts significantly higher than CH, (see 6.3))."," However, the $_3$ OH yield is significantly higher than $_4$ (see \ref{discon}) )."
+ The solid state C2H5OH yields are <20%.., The solid state $_2$ $_5$ OH yields are $\leq$.
+ Previously. C2HsOH and CH;CHO. were shown to form in interstellar ice analogues by photolysis of C4H*:H?O mixtures (Mooreetal.2001:Wu2002).," Previously, $_2$ $_5$ OH and $_3$ CHO were shown to form in interstellar ice analogues by photolysis of $_2$ $_2$ $_2$ O mixtures \citep{moore2001,wu2002}."
+. Since in such experiments both OH and H fragments are present with excess energy. it is difficult to disentangle potential pure hydrogenation reactions and reactions involving OH radicals.," Since in such experiments both OH and H fragments are present with excess energy, it is difficult to disentangle potential pure hydrogenation reactions and reactions involving OH radicals."
+ Indeed. Moore&Hudson(2005) explain. formation of C;H3OH and CH3CHO by reactions of C4Hs and C? Hs with OH. respectively.," Indeed, \citet{moore2005} explain formation of $_2$ $_5$ OH and $_3$ CHO by reactions of $_2$ $_5$ and $_2$ $_3$ with OH, respectively."
+ In this paper we focus on the reactions with thermal H-atoms only., In this paper we focus on the reactions with thermal H-atoms only.
+ Since H:CO is known to react with H-atoms to CH;OH (Watanabeetal..2004:Hidaka2004) itis thus likely that the next more complex aldehyde. acetaldehyde (CH;CHO). will form ethanol (C>H;5OH).," Since $_2$ CO is known to react with H-atoms to $_3$ OH \citep{watanabe2004,hidaka2004} it is thus likely that the next more complex aldehyde, acetaldehyde $_3$ CHO), will form ethanol $_2$ $_5$ OH)."
+ In Figure 12. the relative heats of formation at 0 K are shown (Wibergetal..199]:Cox 2007).," In Figure \ref{energy} the relative heats of formation at 0 K are shown \citep{wiberg1991,cox1989,gurvich1989,frenkel1994,matus2007}."
+. The exothermicity of CH; CHOH formation is higher than that of CH:CH?O. but which of the species is more likely formed depends on the reaction barriers.," The exothermicity of $_3$ CHOH formation is higher than that of $_3$ $_2$ O, but which of the species is more likely formed depends on the reaction barriers."
+ Subsequent formation of C+Hs5OH is likely fast. because reactions of radicals with H- commonly have no activation barriers.," Subsequent formation of $_2$ $_5$ OH is likely fast, because reactions of radicals with H-atoms commonly have no activation barriers."
+ As described in ?? only a fraction of CH;CHO ts converted to C4 H&SOH. and a larger fraction leads to CH4y. H:CO. and CH3OH formation.," As described in \ref{sec_rate_ch3cho} only a fraction of $_3$ CHO is converted to $_2$ $_5$ OH, and a larger fraction leads to $_4$, $_2$ CO, and $_3$ OH formation."
+ For hydrogenation of CH:CHO a C=O bond is converted to a C-O bond instead of breaking a C-C bond., For hydrogenation of $_3$ CHO a C=O bond is converted to a C–O bond instead of breaking a C–C bond.
+ Since the C=O bond is intrinsically stronger it is likely that the entrance channel to hydrogenation is higher in energy compared to dissociation., Since the C=O bond is intrinsically stronger it is likely that the entrance channel to hydrogenation is higher in energy compared to dissociation.
+ Thus H-atoms can break the C-C bond as well as the C=O bond to form CH;. H:CO and CH;OH or C:HsOH in ices as prepared here.," Thus H-atoms can break the C–C bond as well as the C=O bond to form $_4$, $_2$ CO and $_3$ OH or $_2$ $_5$ OH in ices as prepared here."
+ As shown in Fig., As shown in Fig.
+" 12. the formation of CH,+HCO is more exothermic than that for CH;+H2CO.", \ref{energy} the formation of $_4$$+$ HCO is more exothermic than that for $_3$$+$ $_2$ CO.
+" Furthermore. the energy released in this step is higher than the binding energy of CH, to the surface. which is ~700 K (0.06 eV)."," Furthermore, the energy released in this step is higher than the binding energy of $_4$ to the surface, which is $\sim$ 700 K (0.06 eV)."
+" This likely explains why the ¥(CH4) is lower than Y(CH:OH-«H:CO.. because the formation energy is sufficient for CH, desorption."," This likely explains why the $Y$ $_4$ ) is lower than $Y$ $_3$ $+$ $_2$ CO), because the formation energy is sufficient for $_4$ desorption."
+ The energy released during the formation of H:CO. CH:OH and C;H5OH ts even higher and may also cause a fraction of the molecules to desorb.," The energy released during the formation of $_2$ CO, $_3$ OH and $_2$ $_5$ OH is even higher and may also cause a fraction of the molecules to desorb."
+ Our experiments show that CO» reaction rates with H-atoms are very low. making it an implausible route for HCOOH formation.," Our experiments show that $_2$ reaction rates with H-atoms are very low, making it an implausible route for HCOOH formation."
+ A number of other HCOOH formation routes are possible (seee.g..Milligan&Jacox.1971;HudsonMoore.1999;Keane. 2001).. from either HCO+OH — HCOOH or HCO+0 — HCOO+H HCOOH.," A number of other HCOOH formation routes are possible \citep[see
+e.g.,][]{milligan1971,hudson1999,keane2001}, from either $+$ OH $\rightarrow$ HCOOH or $+$ O $\rightarrow$ $+$ H $\rightarrow$ HCOOH."
+ In addition. experiments suggest that under specific catalytic conditions CO» can react to form HCOOH (Ogoetal..2006) but this requires catalytic surface sites. Le. CO» directly attached to à. silicate. or metallic grain site.," In addition, experiments suggest that under specific catalytic conditions $_2$ can react to form HCOOH \citep{ogo2006} but this requires catalytic surface sites, i.e., $_2$ directly attached to a silicate or metallic grain site."
+ Such a situation is less likely in dense interstellar clouds where thick ice layers have already formed and cover any potential catalytic sites., Such a situation is less likely in dense interstellar clouds where thick ice layers have already formed and cover any potential catalytic sites.
+ In conclusion. under astrophysically relevant conditions solid CO» in bulk ice is à very stable molecule that is not likely to react with H-atoms.," In conclusion, under astrophysically relevant conditions solid $_2$ in bulk ice is a very stable molecule that is not likely to react with H-atoms."
+ Similar to CO». reaction rates of HCOOH with H-atoms are below the detection limit in our experiment.," Similar to $_2$, reaction rates of HCOOH with H-atoms are below the detection limit in our experiment."
+ Formation of the so far undetected interstellar species CH2(OHs in this way thus seems unlikely., Formation of the so far undetected interstellar species $_2$ $_2$ in this way thus seems unlikely.
+ Unless other formation mechanisms are found an observational search for this species based upon solid state astrochemical arguments is not warranted., Unless other formation mechanisms are found an observational search for this species based upon solid state astrochemical arguments is not warranted.
+ We conclude that CO:. HCOOH and CH+(OH): do not appear to be related through successive hydrogenation in interstellar ice analogues under the conditions as used in the present study.," We conclude that $_2$, HCOOH and $_2$ $_2$ do not appear to be related through successive hydrogenation in interstellar ice analogues under the conditions as used in the present study."
+ In contrast to CO» and HCOOH. CO does react with H-," In contrast to $_2$ and HCOOH, CO does react with H-atoms."
+ The reaction rates of CO in CO:CO» mixtures are very similar to those found by Fuchsetal.(2007) for pure CO tices., The reaction rates of CO in $_2$ mixtures are very similar to those found by \citet{fuchs2007} for pure CO ices.
+ CO hydrogenation in interstellar ices will thus not be strongly affected by the presence of CO» in the ice., CO hydrogenation in interstellar ices will thus not be strongly affected by the presence of $_2$ in the ice.
+ It is likely that the reaction rate is the same for H+CO independent of the CO concentration and that of other species in apolar interstellar ices., It is likely that the reaction rate is the same for $+$ CO independent of the CO concentration and that of other species in apolar interstellar ices.
+to that of the galaxwv aud then it is applicable for a stualler radius.,to that of the galaxy and then it is applicable for a smaller radius.
+ Ou the other laud. theChandra analysis yields Mg.(4100kpe)=(2.5£0.2)ο”. where we adopted ithe mean metallicity measured witlin GOOTs) kpe. 0.76 solu.," On the other hand, the analysis yields $M_{\rm Fe}(<400\,h_{50}^{-1}{\rm
+kpc})=(2.5\pm0.2)\times10^{10}h_{50}^{-2.5}{\rm M_{\sun}}$ where we adopted the mean metallicity measured within $600\,h_{50}^{-1}$ kpc, 0.76 solar."
+ Because this is by about lugher than the above calculation. there iav be au abuudance variation from cluster to cluster. as suggested from the recent results of metallicity maps: the Inhomoecneous nietal distributious were fouud iu the core reeious of some uearby clusters such as Perseus Fabian.&Sanders2002) and A2199 (Johustoueetal. 2002).," Because this is by about higher than the above calculation, there may be an abundance variation from cluster to cluster, as suggested from the recent results of metallicity maps: the inhomogeneous metal distributions were found in the core regions of some nearby clusters such as Perseus \citep{Schmidt_etal_2002} and A2199 \citep{Johnstone_etal_2002}."
+. However. since the iron mass of CLOO2L)17 is consistent with the Mg.Ly relation within the lo errors. the variation is not statisticallygaso significant.," However, since the iron mass of CL0024+17 is consistent with the $M_{\rm Fe}-L_{\rm V,E+SO}$ relation within the $\sigma$ errors, the variation is not statistically significant."
+ To confirm this. we need future observations of the metallicity map with higher seusitivities.," To confirm this, we need future observations of the metallicity map with higher sensitivities."
+ The observed high metallicity can be explained bx metal ejection from the elliptical galaxies., The observed high metallicity can be explained by metal ejection from the elliptical galaxies.
+" Since the spiral ealaxy fraction. fy,=(0.10. is not different from other clusters with a simular mass (σαetal.1997).. the above result mdicates that very effective galaxy formation had occurred in the central region of CLO021|17."," Since the spiral galaxy fraction, $f_{\rm
+sp}=0.40$, is not different from other clusters with a similar mass \citep{Smail_etal_1997}, the above result indicates that very effective galaxy formation had occurred in the central region of CL0024+17."
+ Besides. the optical observations showed that the spatial distribution of ealaxies is ceutrally concentrated aud las a very compact core (Sunudletal.1996:Schueider1986).," Besides, the optical observations showed that the spatial distribution of galaxies is centrally concentrated and has a very compact core \citep{Smail_etal_1996, Schneider_etal_1986}."
+.. Then because of the high concentration of bright elliptical ealaxies the inuetal will also follow the conceutrated distribution., Then because of the high concentration of bright elliptical galaxies the metal will also follow the concentrated distribution.
+ As mentioned above. there nav also be a conrplex metal distribution in the cluster core.," As mentioned above, there may also be a complex metal distribution in the cluster core."
+ However due to the poor photon statistics. we are not able to derive the 2-dimensional abundauce map.," However due to the poor photon statistics, we are not able to derive the 2-dimensional abundance map."
+ For this reason we studied the radiallv-averaged abundance profile iu refsubsecitprof.., For this reason we studied the radially-averaged abundance profile in \\ref{subsec:tprof}.
+ Our result shown in Figure 3bb indicates that the metal abuudince is as high as ~1 solar at the innermost region and then not iu conflict with au abuudance profile in which Zg.~1l within the central rXLO” region aud Zp.~0.3 at the outer region.," Our result shown in Figure \ref{fig:tprof}b b indicates that the metal abundance is as high as $\sim 1$ solar at the innermost region and then not in conflict with an abundance profile in which $Z_{\rm Fe}
+\sim 1$ within the central $r\lesssim10\arcsec$ region and $Z_{\rm Fe}
+\sim 0.3$ at the outer region."
+ In order to further constrain the spatial distribution aud the history of the metal production. better photon statistics for the A-rav spectrum are required.," In order to further constrain the spatial distribution and the history of the metal production, better photon statistics for the X-ray spectrum are required."
+ We found from theChandre observation that tle X-ray surface brightness distribution is mainly described by a superposition of two extended componcuts well fitted with spherical profiles., We found from the observation that the X-ray surface brightness distribution is mainly described by a superposition of two extended components well fitted with spherical $\beta$ -profiles.
+ Furthermore we do not fiud amy sienificaut temperature structure., Furthermore we do not find any significant temperature structure.
+ Thus it is likely that the eas is relaxed in the cluster potential aud. therefore. that hydrostatic equilibrium will be a eood approximation in the X-ray lass estimation.," Thus it is likely that the gas is relaxed in the cluster potential and, therefore, that hydrostatic equilibrium will be a good approximation in the X-ray mass estimation."
+ We further discuss auy possible cause of the mass discrepancy below., We further discuss any possible cause of the mass discrepancy below.
+ As shown in the previous section. we can consider tha the N-xav temperature represeuts the virial temperature of the cluster component since the gas mass fraction of the cluster component within the virial radius is consisten with the universal birvou fraction obtained by WALAP.," As shown in the previous section, we can consider that the X-ray temperature represents the virial temperature of the cluster component since the gas mass fraction of the cluster component within the virial radius is consistent with the universal baryon fraction obtained by WMAP."
+ Therefore. the mass discrepancy reported in the previous section is telling us that there is a lack of our curren understanding for the nature of the cluster ceutral region.," Therefore, the mass discrepancy reported in the previous section is telling us that there is a lack of our current understanding for the nature of the cluster central region."
+ Iu some uearby clusters. particularly XD clusters (Forman&Jones1990).. an additional cool comiponeu is required to explain the XN-rav cussion of the central ~LOOhey kpe regions (ATaksishimaetal.2001).," In some nearby clusters, particularly XD clusters \citep{Forman_Jones_1990}, an additional cool component is required to explain the X-ray emission of the central $\sim100~h_{50}^{-1}$ kpc regions \citep{Makishima_etal_2001}."
+. Ou the contrary Makishiniactal.(2001) also suggested that non-ND clusters appear to be isothermal with little metallicity eradicut toward the ceuter., On the contrary \cite{Makishima_etal_2001} also suggested that non-XD clusters appear to be isothermal with little metallicity gradient toward the center.
+ Taking into account the fact that the current target cluster is a non-ND cluster aud that the total A-ray cuuission is not domüuated by the central ( enussion. our results shown iu refsec:spec seein to be more consistent with their picture of non-XD systems.," Taking into account the fact that the current target cluster is a non-XD cluster and that the total X-ray emission is not dominated by the central G1 emission, our results shown in \\ref{sec:spec} seem to be more consistent with their picture of non-XD systems."
+ Towever if the ICAL of CLOO21)17 is ina two-phase state and the thermal pressure of the eas imn cach phase balances against gravitation. it would ercatlv increase the X-ray mass estimation.," However if the ICM of CL0024+17 is in a two-phase state and the thermal pressure of the gas in each phase balances against gravitation, it would greatly increase the X-ray mass estimation."
+ Thus we attempted o fit the spectra for both (1) the cluster region outside 100 kpe from Cl. ie. 0.26«cr«1.5 and (2) the Cil region. r«0.26 with a two-component NEIAL model.," Thus we attempted to fit the spectra for both (1) the cluster region outside 100 kpc from G1, i.e. $0\arcmin.26<r<1\arcmin.5$ and (2) the G1 region, $r<0\arcmin.26$ with a two-component MEKAL model."
+" Because we are not able to constrain the two temperatures simultaneously uuder the current photon statistics, we fixed the temperature and the metallicity for the cool phase respectively at 1 keV aud 1 solu while the spectral jormalization was adjustable."," Because we are not able to constrain the two temperatures simultaneously under the current photon statistics, we fixed the temperature and the metallicity for the cool phase respectively at 1 keV and 1 solar while the spectral normalization was adjustable."
+ For the region (1). we fouud hat the ux of the cool phase is oulv about of the hot phase and the temperature of the hot-phase is 5.5 (L1 7.2) keV. which is consistent with the result of the sinele-xiase model shown in Table P. within errors.," For the region (1), we found that the flux of the cool phase is only about of the hot phase and the temperature of the hot-phase is 5.5 (4.4 – 7.2) keV, which is consistent with the result of the single-phase model shown in Table \ref{tab:spec} within errors."
+ Thus the spectrum is well represented by the single-phase model., Thus the spectrum is well represented by the single-phase model.
+ This is also consistent with the results of nearby won-ND clusters., This is also consistent with the results of nearby non-XD clusters.
+ For the region (2). because the radiative cooling time scale of the C1 emission is shorter than the age of the Universe (Sec refsubsec:xrayvpropertics)}. the gas nieht be in a two-phase state.," For the region (2), because the radiative cooling time scale of the G1 emission is shorter than the age of the Universe (See \\ref{subsec:xrayproperties}) ), the gas might be in a two-phase state."
+ Towever. from the two-phase model fitting. we found that the wpper limit of the cool-phase fux is less than of the hot-plase and the hot-phase temperature. [2 (3.5 5.3) keV. is again consistent with that obtained for the single-phase model.," However, from the two-phase model fitting, we found that the upper limit of the cool-phase flux is less than of the hot-phase and the hot-phase temperature, 4.2 (3.5 – 5.3) keV, is again consistent with that obtained for the single-phase model."
+ Thus we consider that it is unlikely that the Gl eas is ina two-phase state aud that the virial teiiperatfure is much higher than that derived from the single-phase model., Thus we consider that it is unlikely that the G1 gas is in a two-phase state and that the virial temperature is much higher than that derived from the single-phase model.
+ Czoskeetal.(2001.2002) measured the redshift distribution of galaxies in the direction of CLOO21117 and revealed the presence of foreground. aud. backerouud eroups of ealaxies that align along the line of sight.," \cite{Czoske_etal_2001, Czoske_etal_2002} measured the redshift distribution of galaxies in the direction of CL0024+17 and revealed the presence of foreground and background groups of galaxies that align along the line of sight."
+" Czoske""metal.(2002) suggested from their results that there was a high speed (3000 km/s) collision between the CL0021117 cluster aud a secoud cluster with a iiass about half that of the main cluster 3 Gyr ago.", \cite{Czoske_etal_2002} suggested from their results that there was a high speed $\sim 3000$ km/s) collision between the CL0024+17 cluster and a second cluster with a mass about half that of the main cluster $\sim3$ Gyr ago.
+ They suggested that taking iuto account the projection effect of the secoud cluster lass may solve the mass discrepancy problem., They suggested that taking into account the projection effect of the second cluster mass may solve the mass discrepancy problem.
+ They also mentioned that diving the collision the N-rav eas would be liehly disturbed due to propagation of shock waves (es.Takizawa1999) and hydrostatic equilibrimu. of the eas component will break down. however. after several Cars the gas will settle down to an equilibrium state.," They also mentioned that during the collision the X-ray gas would be highly disturbed due to propagation of shock waves \citep[e.g.][]{TAKIZAWA.1999} and hydrostatic equilibrium of the gas component will break down, however, after several Gyrs the gas will settle down to an equilibrium state."
+ Rocttineerctal.(1996) noted that. based on nunercal simulations. the temperature structure is oue of the strougest indicator of recent merger activity. however. we do not fud any substantial temperature structure in the observational data as shown in Figure 3((a).," \cite{Roettinger_etal_1996} noted that, based on numerical simulations, the temperature structure is one of the strongest indicator of recent merger activity, however, we do not find any substantial temperature structure in the observational data as shown in Figure \ref{fig:tprof}( (a)."
+ We also found from the spectral analvsis of the Cl peak that the teiiperature is still hieher than 2.6 keV at the central rocLY vyosion., We also found from the spectral analysis of the G1 peak that the temperature is still higher than 2.6 keV at the central $r<4\arcsec$ region.
+ Thus our results show that the cluster, Thus our results show that the cluster
+"For more general cases, the simple solution of the second order equation (20)) has to be used.","For more general cases, the simple solution of the second order equation \ref{2ndeq}) ) has to be used."
+ Most critical of course are the values ofNÀ and N.2 which take large values in a narrow region surrounding the central core in the helium and more advanced evolutionary stages.," Most critical of course are the values of$N^2_{\Omega}$ and $N^2_{\mu}$, which take large values in a narrow region surrounding the central core in the helium and more advanced evolutionary stages."
+" In order to examine quantitatively the importance of the GSF instability, we calculate the evolution all the way from the Main Sequence to the Si burning stage of a 20 Mo star with an initial rotation velocity of 150 km s-! with a metallicity Z—0.002 typical of the SMC composition."," In order to examine quantitatively the importance of the GSF instability, we calculate the evolution all the way from the Main Sequence to the Si burning stage of a 20 $_{\odot}$ star with an initial rotation velocity of 150 km $^{-1}$ with a metallicity Z=0.002 typical of the SMC composition."
+" We make the choice of this composition, because the internal Q-gradients are steeper at lower Z (Maeder&Meynet 2001)), which would favor the GSF instability."," We make the choice of this composition, because the internal $\Omega$ –gradients are steeper at lower $Z$ \cite{MMVII}) ), which would favor the GSF instability."
+ Some data for another 20 Mo model with an initial rotation of 300 km s! are also given., Some data for another 20 $_{\odot}$ model with an initial rotation of 300 km $^{-1}$ are also given.
+ Equation (20)) was used to determine the occurence of the GSF instability and the value of Dagr., Equation \ref{2ndeq}) ) was used to determine the occurence of the GSF instability and the value of $D_{\mathrm{GSF}}$.
+ The above expression (10)) for Dy is used., The above expression \ref{dhm}) ) for $D_{\mathrm{h}}$ is used.
+ The nuclear network in the advanced phases is the same as in previous models (Hirschietal. 2004))., The nuclear network in the advanced phases is the same as in previous models \cite{HMMXII}) ).
+ Figure 1 shows in four panels the main parameters during the first part of the phase of central He-burning., Figure \ref{MHe} shows in four panels the main parameters during the first part of the phase of central He–burning.
+ We first notice in panel d) the building of a Q-gradient at the edge of the convective core with a difference of Q by about a factor of 20., We first notice in panel d) the building of a $\Omega$ –gradient at the edge of the convective core with a difference of $\Omega$ by about a factor of 20.
+ This makes dlnQ/dInr+2«0 in most of the region between the edge of the convective core at 2.9 Ma and the convective H-burning shell at 5.3 Mo as shown in panel b).," This makes ${\rm d\,ln} \,\Omega/ {\rm d\,ln} \,r +2 <0$ in most of the region between the edge of the convective core at 2.9 $_{\odot}$ and the convective H–burning shell at 5.3 $_{\odot}$ as shown in panel b)."
+ However N@ remains negligiblewith respect to N2 and NPt, However $N^2_{\Omega}$ remains negligiblewith respect to $N^2_T$ and $N^2_{\mu}$.
+" In order to understand why, we need to look back at Eq. (15)):"," In order to understand why, we need to look back at Eq. \ref{n2o}) ):"
+ NZ=20?(24- dn2).," $N^2_{\Omega}=2 \overline{\Omega}^2 ( 2 + 
+\frac{{\rm d\,ln}\,\overline{\Omega}}{{\rm d\,ln}\,r})$ ."
+ The value of 9? in the star is too small to allow a significant value of N2., The value of $\Omega^2$ in the star is too small to allow a significant value of $N^2_{\Omega}$ .
+ This means in fact that the centrifugal force, This means in fact that the centrifugal force
+As discussed above. if only the first ten percent or so of matter in a mass transfer event can be accreted (e.g. Packet 1981: Dewi 2006: but see also Petrovic et al.,"As discussed above, if only the first ten percent or so of matter in a mass transfer event can be accreted (e.g. Packet 1981; Dewi 2006; but see also Petrovic et al."
+ 2005. in which mass transfer is linked to rotation in such a way that an initial mass transfer even can be non-conservative but a later event more nearly conservative) then the potential metallicity gain-like effects of accretion would be basically negligible.," 2005, in which mass transfer is linked to rotation in such a way that an initial mass transfer event can be non-conservative but a later event more nearly conservative) then the potential metallicity gain-like effects of accretion would be basically negligible."
+ However. if some proportion of systems do aecrete a large amount of helium-enriched matter. then this could produce potentially observable effects. at least in individua systems.," However, if some proportion of systems do accrete a large amount of helium-enriched matter, then this could produce potentially observable effects, at least in individual systems."
+ Thus observation of these effects could potentially be used to provide some constraint on the amount of matter accreted in a binary., Thus observation of these effects could potentially be used to provide some constraint on the amount of matter accreted in a binary.
+ However. as noted previously. it is probable that the amoun of matter which can be accreted is variable and affected by a wide range of other parameters. not least rotation (Langer et al.," However, as noted previously, it is probable that the amount of matter which can be accreted is variable and affected by a wide range of other parameters, not least rotation (Langer et al."
+ 2003)., 2003).
+ In our previous work modelling binary populations 6with slightly less detailed models: see e.g. Dray Tout 2006). based on the 20 known WR binaries with measured masses in the catalogue of van der Hucht (2001) we found that no particular accretion scenario is favoured for all stars — most systems which are easy to fit if we assume conservative evolution are also easy to fit using non-conservative evolution but different starting parameters.," In our previous work modelling binary populations (with slightly less detailed models; see e.g. Dray Tout 2006), based on the 20 known WR binaries with measured masses in the catalogue of van der Hucht (2001) we found that no particular accretion scenario is favoured for all stars – most systems which are easy to fit if we assume conservative evolution are also easy to fit using non-conservative evolution but different starting parameters."
+ Some systems require nearly-conservative evolution (e.g. Dray 2006) and some require non-conservative evolution., Some systems require nearly-conservative evolution (e.g. Dray 2006) and some require non-conservative evolution.
+ There is a further group of binaries whose period is too small to tit well cand which are therefore presumably post-common envelope svstems. whose evolution we have not followed in detail due to the large uncertainty surrounding common envelope evolution) and a couple of unusual systems whose origin may be dynamical.," There is a further group of binaries whose period is too small to fit well (and which are therefore presumably post-common envelope systems, whose evolution we have not followed in detail due to the large uncertainty surrounding common envelope evolution) and a couple of unusual systems whose origin may be dynamical."
+ Some dynamically-formed systems are likely because some stars have exchanged companions in dense environments in the past CVanbeveren 2005)., Some dynamically-formed systems are likely because some stars have exchanged companions in dense environments in the past (Vanbeveren 2005).
+ This may also be the case in those systems which contain two vary massive stars. such as WR20a and WR47. although these are expected if the initial binary distribution is skewed towards twin systems in which both stars start with similar masses (Pinsonneault Stanek 2006). perhaps as a natural result of the mode of massive star formation (Bonnell Bate 2006).," This may also be the case in those systems which contain two vary massive stars, such as WR20a and WR47, although these are expected if the initial binary distribution is skewed towards twin systems in which both stars start with similar masses (Pinsonneault Stanek 2006), perhaps as a natural result of the mode of massive star formation (Bonnell Bate 2006)."
+ Investigations of whole-»opulation properties with these models (Dray et al., Investigations of whole-population properties with these models (Dray et al.
+ 2005: Dray 2006) again suggested a picture in which some mass transfer is conservative and some not., 2005; Dray 2006) again suggested a picture in which some mass transfer is conservative and some not.
+ In this rather complex scenario. he uncertainty associated with mass transfer is so much greater han that associated with the precise effects of rejuvenation that not only is it difficult to say anything concrete about the latter Tom comparison of whole-population properties with models. but constraints on the mass accreted in an individual binary. if not coupled to a detailed understanding of that binary previous evolution. may say little about the mass transfer situation sin general other that providing a rather weak limit.," In this rather complex scenario, the uncertainty associated with mass transfer is so much greater than that associated with the precise effects of rejuvenation that not only is it difficult to say anything concrete about the latter from comparison of whole-population properties with models, but constraints on the mass accreted in an individual binary, if not coupled to a detailed understanding of that binary's previous evolution, may say little about the mass transfer situation in general other that providing a rather weak limit."
+ It may still be possible to observe the effects of enhanced-metallicity rejuvenation in individual stars and groups of stars., It may still be possible to observe the effects of enhanced-metallicity rejuvenation in individual stars and groups of stars.
+ In particular. we expect from our models that these effects would be most visible in the time period after the SN explosion of the primary in à system which has been undergoing mass transfer (if the SN order is not reversed — see e.g. Pols 1994).," In particular, we expect from our models that these effects would be most visible in the time period after the SN explosion of the primary in a system which has been undergoing mass transfer (if the SN order is not reversed – see e.g. Pols 1994)."
+ If the S does not unbind the system. then the result would be a massive star — compact object binary.," If the SN does not unbind the system, then the result would be a massive star – compact object binary."
+ Cyg X—3 is an example of such a system in which there is a WR star: however. it is very unusua (Lommen et al.," Cyg X–3 is an example of such a system in which there is a WR star; however, it is very unusual (Lommen et al."
+ 2005)., 2005).
+ The most common fate. as discussed in Dray et al. (," The most common fate, as discussed in Dray et al. ("
+2005) is that the system is split by the SN Kick. and the potentially-rejuvenated component becomes a high-velocity runaway star.,"2005) is that the system is split by the SN kick, and the potentially-rejuvenated component becomes a high-velocity runaway star."
+ Many runaway O stars have been observed to have enhanced surface helium abundances (Blaauw 1993). as found in our models.," Many runaway O stars have been observed to have enhanced surface helium abundances (Blaauw 1993), as found in our models."
+ A notable subgroup of the WR population is the WNS stars. which differ significantly from normal WR stars (Marchenko et al.," A notable subgroup of the WR population is the WN8 stars, which differ significantly from normal WR stars (Marchenko et al."
+ 1998)., 1998).
+ Many of these stars are measurable runaways or appear to have moved a significant distance from their place of birth and very few are in binaries., Many of these stars are measurable runaways or appear to have moved a significant distance from their place of birth and very few are in binaries.
+ Whilst it has been suggested that they could be the result of massive star-compact object mergers (i.e. a type of Thorne-Zyytkow object. Vanbeveren et al.," Whilst it has been suggested that they could be the result of massive star–compact object mergers (i.e. a type of Thorne-Żyytkow object, Vanbeveren et al."
+ 1998: Cherepashehuk Moffat 199-4). it is quite possible that they are just the unusual. WR stars which result from secondary evolution in an accreting binary.," 1998; Cherepashchuk Moffat 1994), it is quite possible that they are just the unusual WR stars which result from secondary evolution in an accreting binary."
+ Tracking of these stars and of runaway O stars back to their places of origin and a comparison of the metallicity of that area with their apparent metallicity from massive star evolutionary tracks and surface compositions could potentially give a limit on how important enhanced-metallicity rejuvenation is. and whether it is an effect that needs to be included in population synthesis models or not.," Tracking of these stars and of runaway O stars back to their places of origin and a comparison of the metallicity of that area with their apparent metallicity from massive star evolutionary tracks and surface compositions could potentially give a limit on how important enhanced-metallicity rejuvenation is, and whether it is an effect that needs to be included in population synthesis models or not."
+ LMD is grateful for support by the Leicester PPARC rolling grant for theoretical astrophysics. and to Peter Eggleton for providing and explaining his code.," LMD is grateful for support by the Leicester PPARC rolling grant for theoretical astrophysics, and to Peter Eggleton for providing and explaining his code."
+ CAT thanks Churchill College for his Fellowship., CAT thanks Churchill College for his Fellowship.
+close to the Magorriau ct al (1998) relation.,close to the Magorrian et al (1998) relation.
+ At that point the cohuun density in to the accretion radius 0r. so the growth is (just) Compton thick.," At that point the column density in to the accretion radius $N_H\sim
+N_T=\sigma_T^{-1}$ , so the growth is (just) Compton thick."
+ The growth of massive black holes is radiatively efficient. highly obscured aud gives rise to much of the NRD.," The growth of massive black holes is radiatively efficient, highly obscured and gives rise to much of the XRB."
+ It is also intimately linked with the erowth of galaxy spheroids. the main evolution of which is terminated by a quasar wind.," It is also intimately linked with the growth of galaxy spheroids, the main evolution of which is terminated by a quasar wind."
+ X-ray observations probe best the uuderlviug obscured nucleus at (rest frame) ebergies of about 30 keV. Indeed N-ravs are the best diagnostic of the black hole accretion history of tle Universe., X-ray observations probe best the underlying obscured nucleus at (rest frame) ebergies of about 30 keV. Indeed X-rays are the best diagnostic of the black hole accretion history of the Universe.
+ The optically bright quasar phase (frou an outside observers point of view) follows over the next few million vears as the accretion disc around the black hole cluptics., The optically bright quasar phase (from an outside observer's point of view) follows over the next few million years as the accretion disc around the black hole empties.
+ The carly phase as the wind clears the gas away can be identified with BAL quasars., The early phase as the wind clears the gas away can be identified with BAL quasars.
+ The ceutral cugine is only revived after the quasar phase if a merecr or other event brines in sufficient low augular momentum gas to fucl it., The central engine is only revived after the quasar phase if a merger or other event brings in sufficient low angular momentum gas to fuel it.
+ The prospects of testing the above scenario auc absorption models of the NRB are close at hand. with Chaudra and NADA," The prospects of testing the above scenario and absorption models of the XRB are close at hand, with Chandra and XMM."
+L They should detect large uuubers of faint. but powerful absorbed sources in the 310 keV band. due to the negative Ix. correction involved (see Fig.," They should detect large numbers of faint, but powerful absorbed sources in the 3–10 keV band, due to the negative K correction involved (see Fig."
+ 2) aud ideutifv them with huninous voune galaxy spheroids., 3) and identify them with luminous FIR/sub-mm--emitting young galaxy spheroids.
+an accretion disk (1.0. it is iudependent of AM).,an accretion disk (i.e. it is independent of $\dot M$ ).
+ Therefore. the above formula is valid provided that the disk is able to transfer van which it may do purely by diffusion if à>LR. ox by warp waves (Lubow Pringle 1995 Papaloizon Liu 1995. Nelson Papaloizou 1998) if αsIIR. aud provided that the disk cau act as au adequate sink of augular moment.," Therefore, the above formula is valid provided that the disk is able to transfer warp, which it may do purely by diffusion if $\alpha 
+> H/R$, or by warp waves (Lubow Pringle 1993, Papaloizou Lin 1995, Nelson Papaloizou 1998) if $\alpha\,<\,H/R$, and provided that the disk can act as an adequate sink of angular momentum."
+" Iu the viscous case. (am ΠΠ) volevaut to AGN disks. this timescale iav be rewritten (Scheucr Feiler 1995) as.(2-11) where X is the surface deusity of the disk at the warp radius Pha,"," In the viscous case, $\alpha\,>\,H/R$ ) relevant to AGN disks, this timescale may be rewritten (Scheuer Feiler 1995) as, where $\Sigma$ is the surface density of the disk at the warp radius $R_{\rm warp}$."
+ For a steady accretion disk far frou the centre. it canbe shown that (Pringle 1981).," For a steady accretion disk far from the centre, it can be shown that (Pringle 1981),."
+for mu a stecodily accreting disk. we may write.ΠΟΤΟ) where the accretion timescale free Is defined as faa AM.," Thus, for a steadily accreting disk, we may write, where the accretion timescale $t_{\rm acc}$ is defined as $t_{\rm acc}\,=\,M/\dot M$ ."
+ Alternatively. for a steady disk. the accretion timescale can be written in terms of the Salpeter time £5. which is the growth timescale for the black hole if it is accreting at a rate corresponding to the laniting Eddinetou luuinosity Lr.," Alternatively, for a steady disk, the accretion timescale can be written in terms of the Salpeter time $t_S$, which is the growth timescale for the black hole if it is accreting at a rate corresponding to the limiting Eddington luminosity $L_E$."
+ Thus.LO where. te=12«105(e/0.3) voars.," Thus, where, $t_S\,=\,1.2\,\times\,10^8\,({\epsilon}/{0.3})$ years."
+ Making use of equations 2-7. 2-8. 2-13 and 2-11. we find that or equivalently. We have shown that our current understanding of the time-cvolution of warps in accretion disks leads to the conclusion that a disk which is misaligned with the spin of a central black hole. aud whose inner regious are therefore aligned with the spin of the hole by the Bardecu-Pettersou effect. brings the spin vector of the hole iuto aligumeut with the spin vector of the outer disk on a timescale which is 1�� shorter than had previously been realised.," Making use of equations 2-7, 2-8, 2-13 and 2-14, we find that or equivalently, We have shown that our current understanding of the time-evolution of warps in accretion disks leads to the conclusion that a disk which is misaligned with the spin of a central black hole, and whose inner regions are therefore aligned with the spin of the hole by the Bardeen-Petterson effect, brings the spin vector of the hole into alignment with the spin vector of the outer disk on a timescale which is much shorter than had previously been realised."
+ Thus. the idea that the maintenance of the jet directions iu extended double radio sources for timescales of up to of order 10° years is due to the fisvheel effect of the coutra spinning black hole is no longer a tenable ouc.," Thus, the idea that the maintenance of the jet directions in extended double radio sources for timescales of up to of order $10^9$ years is due to the `flywheel' effect of the central spinning black hole is no longer a tenable one."
+ This shouk perhaps iu anv case not be too surprising since for some lud of flywheel mechauisin to work oue would normally select as the flywheel an object with as large a moment of inertia as possible., This should perhaps in any case not be too surprising since for some kind of flywheel mechanism to work one would normally select as the flywheel an object with as large a moment of inertia as possible.
+ Taking cognizance of this. the idea that the smallest object in the svstem nüeht be the firwhlee seenis in retrospect counter-intuitive.," Taking cognizance of this, the idea that the smallest object in the system might be the flywheel seems in retrospect counter-intuitive."
+ Thus a more likely identification as the flmwheel might be the accretion disk itself., Thus a more likely identification as the flywheel might be the accretion disk itself.
+ For example. the radius at which black hole an disk angular momenta are equal for the Collin-Soufftin Dumont (1990) AGN models is," For example, the radius at which black hole and disk angular momenta are equal for the Collin-Souffrin Dumont (1990) AGN models is."
+ Iu this picture. the constancy of directionality of the jets is due to the fact that the accretion event was of a single eas vichvals object.," In this picture, the constancy of directionality of the jets is due to the fact that the accretion event was of a single gas rich object."
+ Alternatively. oue might consider using the ost as the flywheel. if the galactic potential is such hat any accreted eas would soon start to orbit iu some xeferred plane.," Alternatively, one might consider using the host galaxy as the flywheel, if the galactic potential is such that any accreted gas would soon start to orbit in some preferred plane."
+ We should note that although. the overall couclusious are unlikely to change. the details of the aligumeut iuescales aud disk warp structure will depeud n the specific accretion disk model applicable i any elven case.," We should note that although the overall conclusions are unlikely to change, the details of the alignment timescales and disk warp structure will depend n the specific accretion disk model applicable in any given case."
+ For illustration. in this paper. we have used the models w Colliu-Souffiu Dinmont (1990).," For illustration, in this paper, we have used the models by Collin-Souffrin Dumont (1990)."
+ ILowever. it should © borne in niuud. that even these models strictly need sole amendment to take account of assvinetrie heating of the disk because it is twisted aud of additional iuterual weating due to the enhanced 7» type of dissipation brought on by the twist.," However, it should be borne in mind, that even these models strictly need some amendment to take account of assymetric heating of the disk because it is twisted and of additional internal heating due to the enhanced $\nu_2$ type of dissipation brought on by the twist."
+ In addition. it may be that the accretion disk structure is quite different from the standard thin disk structure envisaged bv Collin-Soutffin. Dumont (1990).," In addition, it may be that the accretion disk structure is quite different from the standard thin disk structure envisaged by Collin-Souffrin Dumont (1990)."
+ For example. if the flow onto the hole is iu the form of an advection-dominated flow (ADAF). for which a~ΠΠ~1. (Naravan Yi. 1995) then from equation 2.L we see that we expect ReapRs~ld. together with a consequently reduced aliguiucut timescale.," For example, if the flow onto the hole is in the form of an advection-dominated flow (ADAF), for which $\alpha\,\sim\,H/R\,\sim\,1$, (Narayan Yi, 1995) then from equation 2.4 we see that we expect $R_{\rm warp}/R_s\,\sim\,1$, together with a consequently reduced alignment timescale."
+/ All these issues will need to be addressed in further work., All these issues will need to be addressed in further work.
+ From an observational point of view. if the jet direction is controlled solely by the black hole spin. then. because the Bardecu-Petterson radius is typically at radii too stall to be observed directly. we would not necessarily expect to observe a correlation between the angular πο vector of eas in the host galaxy. and the directious of the jets.," From an observational point of view, if the jet direction is controlled solely by the black hole spin, then, because the Bardeen-Petterson radius is typically at radii too small to be observed directly, we would not necessarily expect to observe a correlation between the angular momentum vector of gas in the host galaxy, and the directions of the jets."
+ In this context. we note that in a survey of nearby radio loudcarly type galaxies with HST. van Dokkuu and Fraux (1995) find that the major axes of the dust disks seen in the inner regions (C 250 pc) are aligned perpendicular to the arc second radio structures in the uuclei.," In this context, we note that in a survey of nearby radio loud early type galaxies with HST, van Dokkum and Franx (1995) find that the major axes of the dust disks seen in the inner regions $\simle$ 250 pc) are aligned perpendicular to the arc second radio structures in the nuclei."
+ This night be taken as evidence that the direction of the radio jets iu these objects is determined by the galactic potential and/or by the orbital angular momentum of the eas rich intruder which prestunably produced both the dust disks and the nuclear activity., This might be taken as evidence that the direction of the radio jets in these objects is determined by the galactic potential and/or by the orbital angular momentum of the gas rich intruder which presumably produced both the dust disks and the nuclear activity.
+ We also note that for this to occur. and indeed for jets," We also note that for this to occur, and indeed for jets"
+"(?),, which in turn further complicates the scheme compared to SPH.","\citep{1999A&A...342..179P}, which in turn further complicates the scheme compared to SPH."
+" Nevertheless, we feel that the advantages of the scheme outweigh these disadvantages."," Nevertheless, we feel that the advantages of the scheme outweigh these disadvantages."
+" Our scheme, in principle, permits adaptive particle splitting to increase resolution in the desired regions."," Our scheme, in principle, permits adaptive particle splitting to increase resolution in the desired regions."
+" However, if not done carefully, this may break the regularity of the particle distribution, and therefore introduce substantial errors in flux divergence."," However, if not done carefully, this may break the regularity of the particle distribution, and therefore introduce substantial errors in flux divergence."
+ The magnitude of these errors and their impact on the outcome of the simulation are hard to estimate., The magnitude of these errors and their impact on the outcome of the simulation are hard to estimate.
+" The MHD test problems with initially non-uniform but properly relaxed particle distribution refsect:blast,,?? and ??)) showed excellent results with initially nested and relaxed particle distribution."," The MHD test problems with initially non-uniform but properly relaxed particle distribution \\ref{sect:blast}, \ref{sect:orszag} and \ref{sect:rotor}) ) showed excellent results with initially nested and relaxed particle distribution."
+" In principle, if a group of particles is added in the course of simulation and their neighbours are adjusted to maintain regular distribution, the noise should be small."," In principle, if a group of particles is added in the course of simulation and their neighbours are adjusted to maintain regular distribution, the noise should be small."
+" However, further research is required to discover optimal way for particle refinement."," However, further research is required to discover optimal way for particle refinement."
+ We showed that the application of the weighted meshless scheme to the equations of ideal hydrodynamics is straightforward., We showed that the application of the weighted meshless scheme to the equations of ideal hydrodynamics is straightforward.
+ We also expect that such scheme is computationally slightly more expensive than SPH., We also expect that such scheme is computationally slightly more expensive than SPH.
+" In its optimal form, it requires two loops over neighbours, compared to one in SPH: calculation and limit of the gradients of primitive fluid variables, and the interaction part."," In its optimal form, it requires two loops over neighbours, compared to one in SPH: calculation and limit of the gradients of primitive fluid variables, and the interaction part."
+" This is in addition to Eq.(15)), which, similarly to SPH, is solved iteratively."," This is in addition to \ref{eq:constraint_h}) ), which, similarly to SPH, is solved iteratively."
+" The interaction part of the scheme requires solution of a Riemann problem between a particle and its neighbours; on average, 32 Riemann problems are solved for each particle in 3D. Nevertheless, both HLL or HLLC solvers are only moderately expensive compared to calculation of artificial viscosity and conductivity in SPH."," The interaction part of the scheme requires solution of a Riemann problem between a particle and its neighbours; on average, 32 Riemann problems are solved for each particle in 3D. Nevertheless, both HLL or HLLC solvers are only moderately expensive compared to calculation of artificial viscosity and conductivity in SPH."
+" However, our scheme has higher memory footprint due to storage of gradients in the memory."," However, our scheme has higher memory footprint due to storage of gradients in the memory."
+" Nevertheless, combined with the lower number of neighbours usually uses in 3D SPH"," Nevertheless, combined with the lower number of neighbours usually uses in 3D SPH"
+deusity aud 5;; is the stress.,density and $S_{ij}$ is the stress.
+ Finally. 8 is defined. as the trace of 5;j. We remark that if the constraint equations are satisfied ou an initial time slice X. the evolution equations euarantee that the coustraimts will be satisfied on all subsequent time slices.," Finally, $S$ is defined as the trace of $S_{ij}$, We remark that if the constraint equations are satisfied on an initial time slice $\Sigma$, the evolution equations guarantee that the constraints will be satisfied on all subsequent time slices."
+ Equations (2-6)) through (2-9)) are commouly referred to as the ADM equations (Arnowitt. Deser Alisuer. 1962).," Equations \ref{ham}) ) through \ref{gdot}) ) are commonly referred to as the ADM equations (Arnowitt, Deser Misner, 1962)."
+ Niwunerical implementations of these equations have revealed that their umuerical stability cau be iniproved dramatically by biugius them into a slightly different form., Numerical implementations of these equations have revealed that their numerical stability can be improved dramatically by bringing them into a slightly different form.
+" Oue such modification. now commonly referred to as ""DSSN. i$ based on Shibata Nakamura (1995) aud Banmearte Shapiro (19995)."," One such modification, now commonly referred to as “BSSN”, is based on Shibata Nakamura (1995) and Baumgarte Shapiro (1999b)."
+ This system is widely used. aud dts cuhauced stability propertics have been analyzed by several authors (e.g. Aleubierre ct al.," This system is widely used, and its enhanced stability properties have been analyzed by several authors (e.g. Alcubierre et al.,"
+ 2000: see Knapp. Walker Batuearte 2002 for an clectromagnetic analogy).," 2000; see Knapp, Walker Baumgarte 2002 for an electromagnetic analogy)."
+ Alternatively. several authors have experimented with Lyperbolic foxiuulatious of Einstein's equations (c.g. Anderson York 1999).," Alternatively, several authors have experimented with hyperbolic formulations of Einstein's equations (e.g. Anderson York 1999)."
+ We refer the reader to these papers for further details and references., We refer the reader to these papers for further details and references.
+" We decompose the Faraday tensor F as so that E"" and BY are the electric aud magnetic fields observed by a normal observer a"".", We decompose the Faraday tensor $F^{ab}$ as so that $E^a$ and $B^a$ are the electric and magnetic fields observed by a normal observer $n^a$.
+ Both fields are purely spatial. whereby aud the threc-dineusional Levi-Civita svibol ει 1s defined by Note that e is non-zero oulv for spatial indices. while €ab- Muay be non-vanishing even if one index is timelike (see equation (3-19)) below).," Both fields are purely spatial, whereby and the three-dimensional Levi-Civita symbol $\epsilon_{abc}$ is defined by Note that $\epsilon^{abc}$ is non-zero only for spatial indices, while $\epsilon_{abc}$ may be non-vanishing even if one index is timelike (see equation \ref{epsilon_convert}) ) below)."
+" We also decompose the electromagnetic current four-vector 7"" according to where p. and J"" are the charge deusity and current as observed by a normal observer a"".", We also decompose the electromagnetic current four-vector $\mathcal{J}^a$ according to where $\rho_e$ and $J^a$ are the charge density and current as observed by a normal observer $n^a$.
+" Note that J is purely spatial. J,=0."," Note that $J^a$ is purely spatial, $J^a n_a = 0$."
+ With these definitions. Maxwell's equations. aud where V ds the four-dineusional covariant derivative operator associated with gar. cui be brought iuto the 3|1 forma (soc. 0.8. Thorne MacDonald 1982).," With these definitions, Maxwell's equations, and where $\nabla$ is the four-dimensional covariant derivative operator associated with $g_{ab}$, can be brought into the $3+1$ form (see, e.g., Thorne MacDonald 1982)."
+ The charge couscrvation equation. which is implied by (3-5)). becomes The special relativistic Maxwell's equations cau be recovered very casily by evaluating the above equations for a Miukowski spacetime with 5;;=f; where fy; is the flat spatial metric i an arbitrary coordinate svsteni. a=lN=O0and οἱ=0.," The charge conservation equation, which is implied by \ref{maxwell1}) ), becomes The special relativistic Maxwell's equations can be recovered very easily by evaluating the above equations for a Minkowski spacetime with $\gamma_{ij} = f_{ij}$, where $f_{ij}$ is the flat spatial metric in an arbitrary coordinate system, $\alpha = 1$ , $K = 0$ and $\beta^i
+= 0$."
+" It is convenient to introduce a four-vector potential A,,. which can be decomposed into where A, is purely spatial. ο=0."," It is convenient to introduce a four-vector potential $\mathcal{A}_a$, which can be decomposed into where $A_a$ is purely spatial, $A_a n^a = 0$."
+ DIusertiug (2-3)). this implies while At=0. as for any spatial vector.," Inserting \ref{n}) ), this implies while $A^t = 0$, as for any spatial vector."
+" In terms of the vector potential. the Faraday tensor cau be written Contracting this equation with e""* yields 2 Note that with this identification. the magnetic field B’ automatically satisfies the coustraimt equation (3-9))."," In terms of the vector potential, the Faraday tensor can be written Contracting this equation with $\epsilon^{abc}$ yields or Note that with this identification, the magnetic field $B^i$ automatically satisfies the constraint equation \ref{divB}) )."
+ It is possible. aud often conveuieut (see Paper IT). to re-write Maxwells equations completely im terms of £; and sl; thereby climinating D;.," It is possible, and often convenient (see Paper II), to re-write Maxwell's equations completely in terms of $E_i$ and $A_i$, thereby eliminating $B_i$."
+" Evaluating (3-15)) for the components &=f aud b=¢ with A; aud A,—n0|JA; vields Usine the definition (3-3)). we can rewrite so that With (3-17)). eiiBP cau be expressed iu teris of A; as lusertiug this iuto (3-20)) vields Tn equation (3-8)). the magnetic field D; cau beclinunated similarly: lusertiug this iuto (3-8)) vields Equations (3-22)) and (3-21)) form a system of equations for £ and A; alone."," Evaluating \ref{faraday2}) ) for the components $a
+= t$ and $b = i$ with $\mathcal{A}_i = A_i$ and $\mathcal{A}_t = -
+\alpha \Phi + \beta^i A_i$ yields Using the definition \ref{epsilon}) ), we can rewrite so that With \ref{B}) ), $\epsilon_{ijk} \beta^j B^k$ can be expressed in terms of $A_i$ as Inserting this into \ref{Adot2}) ) yields In equation \ref{Edot}) ), the magnetic field $B_i$ can beeliminated similarly: Inserting this into \ref{Edot}) ) yields Equations \ref{Adot}) ) and \ref{Edot2}) ) form a system of equations for $E^i$ and $A_i$ alone."
+ In the special relativistic, In the special relativistic
+The ADI observing technique was described and its performance using Altair/NIRI at Genmini was presented.,The ADI observing technique was described and its performance using Altair/NIRI at Gemini was presented.
+ It was shown that faint companions can be detected with better 5/N when compared to classical observing techniques lor a wide range of declinations., It was shown that faint companions can be detected with better S/N when compared to classical observing techniques for a wide range of declinations.
+ The ADI technique produces a reference PSF trom the same (areel imaging sequence. removing the need to move to a nearby star for PSF calibration or (ο acquire skv exposures (lor /Z/-band imaging).," The ADI technique produces a reference PSF from the same target imaging sequence, removing the need to move to a nearby star for PSF calibration or to acquire sky exposures (for $H$ -band imaging)."
+ Since (he reference PSF is built using images acquired minutes apart. the reference PSF shows a good quasi-static speckle correlation.," Since the reference PSF is built using images acquired minutes apart, the reference PSF shows a good quasi-static speckle correlation."
+ The stabilitv of the PSF plays a crucial role in ADI as it not only determines the speckle attenuation from (he reference image subtraction but it also determines the regime in which (he noise is attenuated. with increasing observing lime., The stability of the PSF plays a crucial role in ADI as it not only determines the speckle attenuation from the reference image subtraction but it also determines the regime in which the noise is attenuated with increasing observing time.
+ Ht was reported that at Gemini with Altair/NIRI using 30s exposures. the PSF evolves on timescales of minutes and the attenuation by subtraction of a reference image reaches ~2—6 [or short time intervals. achieving better speckle attenuation with better seeing conditions.," It was reported that at Gemini with Altair/NIRI using 30s exposures, the PSF evolves on timescales of $\sim 10-60$ minutes and the attenuation by subtraction of a reference image reaches $\sim 2-6$ for short time intervals, achieving better speckle attenuation with better seeing conditions."
+ The observations of HIP18359. for which a filter change during the sequence reduced significantly the speckle attenuation. underscore the importance of maintaining (he optical setup fixed during the sequence.," The observations of HIP18859, for which a filter change during the sequence reduced significantly the speckle attenuation, underscore the importance of maintaining the optical setup fixed during the sequence."
+" It was shown that the gain in S/N will increasing total observing time for separation greater (han 2"" reaches more than of the optimal case. indicating that the noise is mostly decorrelated between residual images for these separations."," It was shown that the gain in S/N with increasing total observing time for separation greater than $^{\prime \prime}$ reaches more than of the optimal case, indicating that the noise is mostly decorrelated between residual images for these separations."
+ Typical residual speckle decorrelation time is of the order of a few minutes., Typical residual speckle decorrelation time is of the order of a few minutes.
+ The speckle noise residuals decorrelate faster for object. having faster FOV rotation., The speckle noise residuals decorrelate faster for object having faster FOV rotation.
+ In all cases. ADI guarantees a larger gain with longer observation sequences.," In all cases, ADI guarantees a larger gain with longer observation sequences."
+ To our knowledge. this is the first time that such behavior is clearly demonstrated for an acquisition and reduction technique designed [or speckle attenuation.," To our knowledge, this is the first time that such behavior is clearly demonstrated for an acquisition and reduction technique designed for speckle attenuation."
+ The wall raised by quasi-static speckles Chat prevents a gain with longer integration lime for standard observing techniques (Alaroisetal.2003.2005) can thus be removed by ADI.," The wall raised by quasi-static speckles that prevents a gain with longer integration time for standard observing techniques \citep{marois2003,marois2005,masciadri2005} can thus be removed by ADI."
+ Comparison with a classical imaging technique shows that ADI achieves 30 times better speckle attenuation in 30 minutes integration time., Comparison with a classical imaging technique shows that ADI achieves 30 times better speckle attenuation in 30 minutes integration time.
+" The noise attennation obtained on Vega was 100. reaching a contrast of ~20 magnitudes al S"" separation (63 AU)."," The noise attenuation obtained on Vega was 100, reaching a contrast of $\sim$ 20 magnitudes at $^{\prime \prime}$ separation (63 AU)."
+" Observations of the voung stars IIDISSO2S and HD97324D. vielded detection limits in difference of magnitude of 11.1-11.9 at 0.8"". similar to the SDI camera al VLT (Am of ~11 at 0.8). which is an optimized speckle suppression instrument."," Observations of the young stars HD18803 and HD97334B yielded detection limits in difference of magnitude of 11.1-11.9 at $^{\prime
+ \prime}$, similar to the SDI camera at VLT $\Delta m$ of $\sim $ 11 at $^{\prime \prime}$ ), which is an optimized speckle suppression instrument."
+ When combined to substellar models ancl estimated. age for these stars. these observations show that ADI is well suited to search for jovian companions having a mass greater than 1-2 Αρ 50-300 AU awav from nearby voung stars.," When combined to substellar models and estimated age for these stars, these observations show that ADI is well suited to search for jovian companions having a mass greater than 1-2 $_{\rm{Jup}}$ 50-300 AU away from nearby young stars."
+ Finally. ADI could easily and advantageously be combined with SSDI. high-order AO and coronagraphy to improve the detection limits of exoplanets and brown dwarls at all separations.," Finally, ADI could easily and advantageously be combined with SSDI, high-order AO and coronagraphy to improve the detection limits of exoplanets and brown dwarfs at all separations."
+More metal-rich. post-imain sequence stars are expected to lose more mass on the branch compared to less metal-rich stars. ancl this effect is directly. observed in the red-giant. branch huninosity functions of several old open clusters including NGC 188 and NGC 6791 (see Fig.,"More metal-rich post-main sequence stars are expected to lose more mass on the red-giant branch compared to less metal-rich stars, and this effect is directly observed in the red-giant branch luminosity functions of several old open clusters including NGC 188 and NGC 6791 (see Fig."
+ 8 in Nalirai et al., 8 in Kalirai et al.
+ 20072)., 2007a).
+ An important question is the significance of (he mass loss process for different metallicities., An important question is the significance of the mass loss process for different metallicities.
+ Nalirai οἱ al. (, Kalirai et al. (
+2005: 2007a) claimed {ο see a stall dependence of mass loss on metallicity in Hyades ([Fe/Hl] = 40.17) and NGC 2099 (M37: [Fe/I1l] = —0.1).,2005; 2007a) claimed to see a small dependence of mass loss on metallicity in Hyades ([Fe/H] = +0.17) and NGC 2099 (M37; [Fe/H] = $-$ 0.1).
+ For a given fraction of low mass WDs produced from severe giant branch mass loss. one can compare with the field metallicity distiibution to infer a critical abundance above which the late stages of stellar evolution are truncated.," For a given fraction of low mass WDs produced from severe giant branch mass loss, one can compare with the field metallicity distribution to infer a critical abundance above which the late stages of stellar evolution are truncated."
+ Our implied assumption is that all old stars above some threshold truncate their evolution above a critical metallicity and none do below i: (his is certainly a simplification given the stochastic nature of mass loss., Our implied assumption is that all old stars above some threshold truncate their evolution above a critical metallicity and none do below it; this is certainly a simplification given the stochastic nature of mass loss.
+ We also note (hat (his process is expected (o occur only for stars that are old enough (and have low enough turnolf mass) to shed their envelopes prior to experiencing a helium flash., We also note that this process is expected to occur only for stars that are old enough (and have low enough turnoff mass) to shed their envelopes prior to experiencing a helium flash.
+ An intermediate nass metal-rich star would still produce an ordinary C/O WD even with enhanced mass oss., An intermediate mass metal-rich star would still produce an ordinary C/O WD even with enhanced mass loss.
+ As a result. the boundary that we derive is an upper limit. and the process may become important for stars wilh even lower metal abundances.," As a result, the boundary that we derive is an upper limit, and the process may become important for stars with even lower metal abundances."
+ According to the Reid et al. (, According to the Reid et al. (
+2007) analysis. the fraction of [Fe/II] > 0 stars in the solar jeighborhood has been more than in the past 10 Gvirs.,"2007) analysis, the fraction of [Fe/H] $>$ 0 stars in the solar neighborhood has been more than in the past 10 Gyrs."
+ If (he mass loss process was signilicant for |Fe/II] > 0. we would expect to find of the field WD population as single ow nass WDs.," If the mass loss process was significant for [Fe/H] $>$ 0, we would expect to find of the field WD population as single low mass WDs."
+ Since the single low mass WD formation rate among all WDs is around4... (he mass loss is not significant enough for |Fe/II]z- 0 to produce WDs with helium cores (0.45 )).," Since the single low mass WD formation rate among all WDs is around, the mass loss is not significant enough for $>$ 0 to produce WDs with helium cores $<$ 0.45 )."
+ The fraction of metal rich stars with [Fe/H] > +0.2 has been more than in the past LO Gvr. therefore the metallicity effect seems to Kick in for more metal-rich stars.," The fraction of metal rich stars with [Fe/H] $>$ +0.2 has been more than in the past 10 Gyr, therefore the metallicity effect seems to kick in for more metal-rich stars."
+ We find that the mass loss becomes significant enough to produce low mass WDs for |Fe/1l] > +0.3 as the fraction of [Fe/H] > 40.3 stars has been >3% for the past 10 Gyr., We find that the mass loss becomes significant enough to produce low mass WDs for [Fe/H] $>$ +0.3 as the fraction of [Fe/H] $>$ +0.3 stars has been $\geq$ for the past 10 Gyr.
+ A cavead in this analvsis is that the Valenti-Fiseher dataset is biased toward more metal rich stars than the complete volume-Iimited. sample of Reid et al. (, A caveat in this analysis is that the Valenti-Fischer dataset is biased toward more metal rich stars than the complete volume-limited sample of Reid et al. (
+2007: see their Figure 6).,2007; see their Figure 6).
+ Therefore. the actual metallicity limit where (he mass loss process is extrenie enough to create low mass WDs may be lower than |Fe/1l] = 40.3.," Therefore, the actual metallicity limit where the mass loss process is extreme enough to create low mass WDs may be lower than [Fe/H] = +0.3."
+ For example. Allende Prieto et al. (," For example, Allende Prieto et al. ("
+2004) found that among the 1477 thin disk stars in (heir extended sample. the fraction of stars with |[Fe/IHI] > 40.2 and +0.3 is andLs... respectively.,"2004) found that among the 1477 thin disk stars in their extended sample, the fraction of stars with [Fe/H] $>$ +0.2 and +0.3 is and, respectively."
+ Observations of several star clusters with different metallicities can help refine the constraint on metallicity al which the mass loss is significant enough to create low mass, Observations of several star clusters with different metallicities can help refine the constraint on metallicity at which the mass loss is significant enough to create low mass
+ v=—53 | 735.5 , $-$ $^{-1}$ $\sim$ 
+andd luminosities for the clusters as well as field red giants.,d luminosities for the clusters as well as field red giants.
+ loweer metallicity [Fe/H] in the range —0.3 to —0.7 dex., er metallicity [Fe/H] in the range $-$ 0.3 to $-$ 0.7 dex.
+ Also. nearly all the observed field stars seem to be older than NGC 6811 and younger than NGC 6791.," Also, nearly all the observed field stars seem to be older than NGC 6811 and younger than NGC 6791."
+non-virialized. lenses is expected to be smaller than in the virtalized lenses. because they have not vet. collapsed. and hence have larger radii (sce WIxXO2 for more details).,non-virialized lenses is expected to be smaller than in the virialized lenses because they have not yet collapsed and hence have larger radii (see WK02 for more details).
+ We found that the variation in the redshift distribution and the sky density of both lens types with iw depends strongly on the power-sspectrum-normalization approach., We found that the variation in the redshift distribution and the sky density of both lens types with $w$ depends strongly on the power-spectrum-normalization approach.
+ Lf Oy is fixed and ex is normalized to the COBE measurements. there is a significant variation in the abundances with ew.," If $\Omega_0$ is fixed and $\sigma_8$ is normalized to the COBE measurements, there is a significant variation in the abundances with $w$."
+ In particular. the sky density of both virialized lenses and non-virialized. lenses drops by a factor of two from w=1 to we2/3.," In particular, the sky density of both virialized lenses and non-virialized lenses drops by a factor of two from $w = -1$ to $w = -2/3$."
+ phis decline. a result of the significant. decrease in cz with ew. occurs despite the faster formation of structure fore Ll.," This decline, a result of the significant decrease in $\sigma_8$ with $w$, occurs despite the faster formation of structure for $w > -1$ ."
+ IE on the other hand. Qo is allowed to vary with ie such that the CODE normalization matches the cluster- normalization. the redshift distributions and skv density change very little with ew: between w=1 and w=2/3 the sky density of both lens types varies. by just ~ 20%..," If, on the other hand, $\Omega_0$ is allowed to vary with $w$ such that the COBE normalization matches the cluster-abundance normalization, the redshift distributions and sky density change very little with $w$; between $w =-1 $ and $w = -2/3$ the sky density of both lens types varies by just $\sim 20$ ."
+. This insubstantial variation is the result of an increase in Qu with t and a less significant drop in es with was compared to the CODIZ normalization with Qu fixed., This insubstantial variation is the result of an increase in $\Omega_0$ with $w$ and a less significant drop in $\sigma_8$ with $w$ as compared to the COBE normalization with $\Omega_0$ fixed.
+ Obtaining a strong constraint on w from the sky density or redshift distribution of weak lenses therefore appears to be contingent on improved. measurements of Qu fron independent observations., Obtaining a strong constraint on $w$ from the sky density or redshift distribution of weak lenses therefore appears to be contingent on improved measurements of $\Omega_0$ from independent observations.
+ Perhaps more promising is the possibility of utilizing the observed ratio of dark lenses to virialized lenses., Perhaps more promising is the possibility of utilizing the observed ratio of dark lenses to virialized lenses.
+ Unlike measurements of the absolute sky density of weak lenses. the ratio is not very sensitiveto the amount of observational noise in theweak-lensing maps since the abundance of both," Unlike measurements of the absolute sky density of weak lenses, the ratio is not very sensitiveto the amount of observational noise in theweak-lensing maps since the abundance of both"
+where F is the driving flux and g is local eravitational acceleration.,where $\tilde F$ is the driving flux and $\tilde g$ is local gravitational acceleration.
+ The radiative cuerey impinging the unit of surface parallel to the direction to the star is from ecolctrical reasons proportinal to £Rr. where Ris the polar radius. aud £F is the radiative flux at radius r.," The radiative energy impinging the unit of surface parallel to the direction to the star is from geometrical reasons proportinal to $FR/r$, where $R$ is the polar radius, and $F$ is the radiative flux at radius $r$."
+ Assunius that the disk is optically thick (see Sect. A.1.)).," Assuming that the disk is optically thick (see Sect. \ref{apdisopdep}) ),"
+ and all incideu radiation is directed upward. we can roughly estimate Taking 9=GAL/r?. the total disk wind mass-loss rate is then eiven by an inteeral of the miass-loss rate per unit of the disk surface ii between the equatorial radius Rog aud the outer disk radius Rout where factor of 2 in Eq.," and all incident radiation is directed upward, we can roughly estimate Taking $\tilde g=GM/r^2$, the total disk wind mass-loss rate is then given by an integral of the mass-loss rate per unit of the disk surface $\dot m$ between the equatorial radius $R_\text{eq}$ and the outer disk radius $R_\text{out}$ where factor of 2 in Eq."
+ comes from the fact that the wind originates from both sides of the disk., comes from the fact that the wind originates from both sides of the disk.
+ Iusertius the mass flux estimate Eq., Inserting the mass flux estimate Eq.
+ aud we derive where (ising substitution «=r/ R) Asstuning the disk wind is uot viscouslv coupled to the disk. the total augular momentum loss rate via the disk wind is As the disk wind originates mainly from the reeions close to the star (with r/R=10. see Fig. À3)).," and we derive where (using substitution $x=r/R$ ) Assuming the disk wind is not viscously coupled to the disk, the total angular momentum loss rate via the disk wind is As the disk wind originates mainly from the regions close to the star (with $r/R\lesssim10$, see Fig. \ref{palfaobr}) ),"
+ where the azimuthal velocity is roughly equal to the I&epleriau one (see Fie. 1)).," where the azimuthal velocity is roughly equal to the Keplerian one (see Fig. \ref{500_2_a}) ),"
+ we can assume ος2eytre) iu Eq., we can assume $v_\phi\approx v_\text{K}(r)$ in Eq.
+ aud consequently the disk wind angular momentum loss rate is bv a factor of Pr(RouBR) lavecr than the angular iuonentuu loss due fo the CAIS wind launched from equator of hvpothetical critically rotating spherical star with radius A., and consequently the disk wind angular momentum loss rate is by a factor of $P_{\frac{1}{2}}(R_\text{out}/R)$ larger than the angular momentum loss due to the CAK wind launched from equator of hypothetical critically rotating spherical star with radius $R$.
+ A inore detailed calculation (see Appendix Appendix À:)) gives a more complicated form of PiGrour) via Eq., A more detailed calculation (see Appendix \ref{secablaap}) ) gives a more complicated form of $P_\ell(x_\text{out})$ via Eq.
+ which shall be used in Eqs.(," which shall be used in Eqs.,"
+20).. instead of Eq., instead of Eq.
+(21).. For an infinite disk (Roy>xm) we derive from Eq., For an infinite disk $R_\text{out}\rightarrow\infty$ ) we derive from Eq.
+ masini disk wind mass-loss rate aud ΤΠ aueular ποιοπα loss rate as For a typical value of azx0.6 (srti¢ka2006) the maximal disk wind nass-loss rate is relatively low. just about 1/25 of the CAIs stellar wind mass-loss rate.," maximum disk wind mass-loss rate and maximum angular momentum loss rate as For a typical value of $\alpha\approx0.6$ \citep{nlteii} the maximal disk wind mass-loss rate is relatively low, just about $1/25$ of the CAK stellar wind mass-loss rate."
+ The structure of the decretion disk aud the radiativelv daiven wind blowiug frou its surface depends on the value of the angular momentum loss F uccded to keep the stellar rotation at or below the critical rate aud on the magnitude of the radiative force., The structure of the decretion disk and the radiatively driven wind blowing from its surface depends on the value of the angular momentum loss $\dot J$ needed to keep the stellar rotation at or below the critical rate and on the magnitude of the radiative force.
+ If the angular momentum loss is sauall. then the disk could be blown away by the radiative orce already. very close to the star.," If the angular momentum loss is small, then the disk could be blown away by the radiative force already very close to the star."
+ In the opposite case. if he augular moment loss is large. then the mass carried away by the disk wind is ncelieible.," In the opposite case, if the angular momentum loss is large, then the mass carried away by the disk wind is negligible."
+ Tn the iutermediate Case the nias and angular uonientuu will be carried partly by the disk aud partly * the disk wind., In the intermediate case the mass and angular momentum will be carried partly by the disk and partly by the disk wind.
+ When the star has to lose angular nonientuni at a rate J to keep at most the critical rotation. the angular momentum will be carried by the stellar wind GÀ. by the disk wind (Tig). aud by the disk itself GI). For the calculation of total disk inass-loss rate the followiug procedure could be used.," When the star has to lose angular momentum at a rate $\dot J$ to keep at most the critical rotation, the angular momentum will be carried by the stellar wind $\dot J_\text{w}$ ), by the disk wind $\dot
+J_\text{dw}$ ), and by the disk itself $\dot J_{\aalpha})$, For the calculation of total disk mass-loss rate the following procedure could be used."
+ For a given stellar aud line-force parameters (Q aud a) the μααι disk wind augular moimoenutun loss Jasx) corresponding to infinite disk Roar>x can be calculated using Eqs., For a given stellar and line-force parameters $\bar Q$ and $\alpha$ ) the maximum disk wind angular momentum loss $\dot J_\text{dw}(\infty)$ corresponding to infinite disk $R_\text{out}\rightarrow\infty$ can be calculated using Eqs.
+(26).. If the ne aneular ποιοται loss that should be carried away by the disk outflow JAy is lower than the maxim one. οJA.<Jalx). then the disk will be completely ablated by the radiation at the radius Row given (from Eq.," If the net angular momentum loss that should be carried away by the disk outflow $\dot J-\dot J_\text{w}$ is lower than the maximum one, $\dot
+J-\dot J_\text{w}<\dot J_\text{dw}(\infty)$, then the disk will be completely ablated by the radiation at the radius $R_\text{out}$ given (from Eq."
+" for J, =) Iu this case the outer disk radius Rou is equal to the radius above which all material is carried away by the disk wind.", for $\dot J_{\aalpha}=0$ ) In this case the outer disk radius $R_\text{out}$ is equal to the radius above which all material is carried away by the disk wind.
+ The corresponding disk wind mass loss rate AG) is then given by Eqs.(20).," The corresponding disk wind mass loss rate $\dot M_\text{dw}(R_\text{out})$ is then given by Eqs.,."
+.(21)... Note however that the formulae discussed in Sect., Note however that the formulae discussed in Sect.
+ ο. are strictly valid oulv in the optically thick part of the disk (see Sect. A.1.))., \ref{secabla} are strictly valid only in the optically thick part of the disk (see Sect. \ref{apdisopdep}) ).
+" Tt the net augular momentii loss rate JTee is larger than the maxim one, ο)JA.> Fal). then the disk will be oulv partly ablated by the radiation."," If the net angular momentum loss rate $\dot J-\dot J_\text{w}$ is larger than the maximum one, $\dot J-\dot J_\text{w}>\dot J_\text{dw}(\infty)$ , then the disk will be only partly ablated by the radiation."
+ The net, The net
+ooceurs.,occurs.
+ This curve results in solid-body rotation at siuall radi aud a Weplerian-tvpe velocity decrease at large radi., This curve results in solid-body rotation at small radii and a Keplerian-type velocity decrease at large radii.
+ For a flat spira ealaxv icliued to the line-ofseh (such as 1915). the rotation curve can be derived from the variaion of radial velocity along the niajor axis. assunidug that the maxinuun velocity at ασ position. corrected for the galaxy inclination. represents the tangential (1.0. circular) velocity at the corresponding ealactocentric radius.," For a flat spiral galaxy inclined to the line-of-sight (such as 4945), the rotation curve can be derived from the variation of radial velocity along the major axis, assuming that the maximum velocity at any position, corrected for the galaxy inclination, represents the tangential (i.e. circular) velocity at the corresponding galactocentric radius."
+ Tuspection of the position-velocity diagrams along the major axis roefFICCOMP.PVDIAGRAMLIucb) suggests the existence of two ecneral velocity reguues.," Inspection of the position-velocity diagrams along the major axis \\ref{FIG.COMP.PVDIAGRAM}a a,b) suggests the existence of two general velocity regimes."
+ The rapidly rotating central molecular cloud las been discussed alreacky in 33.2., The rapidly rotating central molecular cloud has been discussed already in 3.2.
+ The second system relates to the main disk of the galaxy. viewed at a position angle of ~ 157.," The second system relates to the main disk of the galaxy, viewed at a position angle of $\sim45^{\circ}$ ."
+" It extends out to Ro~100"". with rotational velocities. reaching Vlr (the correction. for au inclination of 787 is oulv 2% and is ignored)."," It extends out to $R\sim400''$, with rotational velocities reaching $V\sim170$ (the correction for an inclination of $78^{\circ}$ is only $2\%$ and is ignored)."
+ This disk is associated with the iuner concentration seen in the distribution of integrated refFIC.AIAPSD))., This disk is associated with the inner concentration seen in the distribution of integrated \\ref{FIG.MAPS}b b).
+" Tn the position-velocity diagriuus the additional coucentrations at R~600"" then correspond to the spiral arms that have already been discussed 33.1).", In the position-velocity diagrams the additional concentrations at $R\sim600''$ then correspond to the spiral arms that have already been discussed 3.4).
+ Because of the high inclination of the galaxy. a line-ofsight at a specific offset along the major axis will also include at ealactoceutric radii lavecr than the radius corresponding to that offset.," Because of the high inclination of the galaxy, a line-of-sight at a specific offset along the major axis will also include at galactocentric radii larger than the radius corresponding to that offset."
+ The lue-ofsight velocitics of the extra components are simaller than the rotational velocity at a specific offset and are responsible for the low-velocity ‘tail’ present at many positions along the major axis., The line-of-sight velocities of the extra components are smaller than the rotational velocity at a specific offset and are responsible for the low-velocity `tail' present at many positions along the major axis.
+ Sinele eaussian profiles were fited. to. individual spectra in the dataset., Single gaussian profiles were fitted to individual spectra in the dataset.
+ To avoid the effects of wings in the spectra. only the top half ofthe profiles with SL το) Shenk /2 was used in the fitting.," To avoid the effects of wings in the spectra, only the top half of the profiles with $S_{\rm L}$ $\geq$ $S_{\rm peak}$ /2 was used in the fitting."
+ Applying auiterative procedure. 11 was fitted to the profile velocities for offsets out to £1507.," Applying aniterative procedure, 1 was fitted to the profile velocities for offsets out to $\pm$ $''$."
+ The parameters listed in Table 5 were obtained for the disk region., The parameters listed in Table \ref{TAB.HI.BRANDT} were obtained for the disk region.
+ They are consistent with the values obtaiuec from a similar fit to 0) data by Dalilem et al. (, They are consistent with the values obtained from a similar fit to $-$ 0) data by Dahlem et al. (
+1993).,1993).
+ Our» values are 3 (for IT1)) aud. from a simular fit to he 1) data. Lowith s1uall formal errors.," Our $n$ values are 3 (for ) and, from a similar fit to the $-$ 1) data, 4, with small formal errors."
+ Our value for 5 is significantly smaller han that determumed by Ables et al. (, Our value for $n$ is significantly smaller than that determined by Ables et al. (
+1987). no.,"1987), $n$ $\sim$ 7."
+ This may be a consequence of the setter angular resolution of our data., This may be a consequence of the better angular resolution of our data.
+ luclusiou of the outermost features m the analysis would have increasedfé: this value is uot well-defined iu the fit to the positiou-velocity diagraii and the siuall uncertainty in Table 5 ouly accounts for the formal errror., Inclusion of the outermost features in the analysis would have increased; this value is not well-defined in the fit to the position-velocity diagram and the small uncertainty in Table \ref{TAB.HI.BRANDT} only accounts for the formal errror.
+ The mass of 11915 can be derived with from the parameters in Table 5 (see e.g. 229 of Draudt 1960: Rogstad et al., The mass of 4945 can be derived with from the parameters in Table \ref{TAB.HI.BRANDT} (see e.g. 29 of Brandt 1960; Rogstad et al.
+ 1967: Dablea ct al., 1967; Dahlem et al.
+ 1993)., 1993).
+ AM is in solar masses. ΕΕ1 5 lis in arcsec. aud D is in Mpc.," $M$ is in solar masses, is in, is in arcsec, and $D$ is in Mpc."
+" For Vis=170 and μμ 3807, AL = 1.1110HAlsol."," For $\vmax = 170$ and = $''$, $M$ = $^{11}$."
+" Reference to the position-velocityv diagrams refFIG.COMDP.PVDIACRAM)) sugecsts that this value uav be a lower Πατ, because the velocities of the outermost features are larger than the values given w the fitted rotation curve."," Reference to the position-velocity diagrams \\ref{FIG.COMP.PVDIAGRAM}) ) suggests that this value may be a lower limit, because the velocities of the outermost features are larger than the values given by the fitted rotation curve."
+" The paucity of at offsets of 104500"" παν sieuifv a lack of eas at the anegential location along the lines-ofsi¢h. in which case he observed velocities are less than the rotation velocities (sce also 33.5.6)."," The paucity of at offsets of $400-500''$ may signify a lack of gas at the tangential location along the lines-of-sight, in which case the observed velocities are less than the rotation velocities (see also 3.5.6)."
+ The total integrated fux deusitv of yvields au mass of 1105 NINE... which is ouly 0.5 of the total mass and considerably less than expected for such ealaxies.," The total integrated flux density of yields an mass of $^{8}$ $_{\odot}$ , which is only $0.5\%$ of the total mass and considerably less than expected for such galaxies."
+ However. the presence of the central absorption has badly affected the iuteeration.," However, the presence of the central absorption has badly affected the integration."
+ Exclusion of the spectzumi observed towards the peak of the ως continua emission (see refFIC.ILCENADSaa) would have duereased the integrated flux densitybv nore thangus.. vieldiug a tota lnass fraction of," Exclusion of the spectrum observed towards the peak of the nuclear continuum emission (see \\ref{FIG.HI.CENABS}a a) would have increased the integrated flux densityby more than, yielding a total mass fraction of."
+ Because of the limited coverage of the CO distribution. he amolecular content of the galaxywas not estimated roni our 1) data.," Because of the limited coverage of the CO distribution, the molecular content of the galaxywas not estimated from our $-$ 1) data."
+" Asstumine an iutegrated CO intensity to IT» column ceusity conversion factor of X = NUT2)/Tew = 110°"" i Dahlen et al. ("," Assuming an integrated CO intensity to $_2$ column density conversion factor of $X$ = $N$ $_2$ $I_{\rm CO}$ = $^{20}$ $^{-1}$ , Dahlem et al. ("
+"1993) derived a molecular eas nius fraction of from their ddata extending out to 2360"".",1993) derived a molecular gas mass fraction of from their data extending out to $''$ .
+ While the X-faetoris, While the $X$ -factoris
+where the sum over jt has to be intended as the sum over all the wave modes present at a given posilion.,where the sum over $\mu$ has to be intended as the sum over all the wave modes present at a given position.
+ In obtaining these relations we explicitly used (he fact that. for Alfvénn waves. £j. in Eq.," In obtaining these relations we explicitly used the fact that, for Alfvénn waves, $F_w$ in Eq."
+ 6 has two contributions: the former is (he normal component of the Povnting vector (BxdT+5Bif)oH/z. While the latter represents the kinetic energy [lux in (rausverse velocity. p/20iu (with du given by Eq. 4)).," \ref{Fw} has two contributions: the former is the normal component of the Poynting vector $(\vec B\times \delta \vec u+\delta\vec B\times\vec u)\times\delta\vec B/4\pi$, while the latter represents the kinetic energy flux in transverse velocity, $\rho/2\delta\vec {u}^2 u$ (with $\delta u$ given by Eq. \ref{disprel}) )."
+" IL the turbulence manifests itself in forms other than resonant Alfénn waves. p, and Ες may differ significantly."," If the turbulence manifests itself in forms other than resonant Alfénn waves, $p_w$ and $F_w$ may differ significantly."
+ We allow the generated waves to have both helicities. hence the sum over the modes should account. in the upstream region. for both forward- and backward-going waves. namely OB.," We allow the generated waves to have both helicities, hence the sum over the modes should account, in the upstream region, for both forward- and backward-going waves, namely $\delta\vec B_{1\pm}$."
+ Each of these modes. however. when advectecl behind the subshock has to split into two waves wilh opposite helicities in order (o satislv the Maxwell equations at the subshock (seee.gMcelxenzie&Westphal1969:Scholer&Beleher 1971).," Each of these modes, however, when advected behind the subshock has to split into two waves with opposite helicities in order to satisfy the Maxwell equations at the subshock \citep[see e.g][]{mck-w69,sb71}."
+. This fact leads to four different downstream modes. which can be described by the introduction of proper transmission and reflection coefficients.," This fact leads to four different downstream modes, which can be described by the introduction of proper transmission and reflection coefficients."
+" In (his scenario a ""reflected"" (""transmitted"") wave is a wave wilh a helicity opposite (equal) to the one of ils upstream counterpart.", In this scenario a “reflected” (“transmitted”) wave is a wave with a helicity opposite (equal) to the one of its upstream counterpart.
+ It is worth stressing that. the fluid being either upstream or downstream. all the modes are actually advected with the [Iuid. independent of their direction of propagation.," It is worth stressing that, the fluid being super-Alfv\'ennic either upstream or downstream, all the modes are actually advected with the fluid, independent of their direction of propagation."
+ These coefficients. 7 and R respectively. were derived by Mclxenzie&Westphal(1969).. and for each upstream helicity 77.=c1 read: M4 =οιis the Alfvénnic Mach number. namely the ratio between {hic and Alfvénn speed. which is of order 100 ancl more for a typical supernova shock.," These coefficients, $\mathcal T$ and $\mathcal R$ respectively, were derived by \cite{mck-w69}, and for each upstream helicity $H_c=\pm 1$ read: Here $M_{A}= u/v_{A}$ is the Alfvénnic Mach number, namely the ratio between fluid and Alfvénn speed, which is of order 100 and more for a typical supernova shock."
+ For each sign of I. we have 0b/dby=T+RBj. aud hence since the subshock can be viewed as à magnetized gas shock. as we already stressed. the jmp conditions found by Vainio&Schlickeiser(1999) for the pressure and temperature hold also in our case and can be written respectively as," For each sign of $H_c$ we have $\delta B_2/\delta B_1=\mathcal
+T+\mathcal R=\Rs$, and hence Since the subshock can be viewed as a magnetized gas shock, as we already stressed, the jump conditions found by \cite{vs99} for the pressure and temperature hold also in our case and can be written respectively as:"
+comparison of the observed aud simulated metallicity distributions for both the elobular clusters aud. halo field stars ds then used to decide which huuinositv fuuctious produce acceptable agreement.,comparison of the observed and simulated metallicity distributions for both the globular clusters and halo field stars is then used to decide which luminosity functions produce acceptable agreement.
+ As in Cotté et al. (, As in Côtté et al. (
+1998). we asstune that the proto-Galactic fragments have equal iierger probabilities: the reader is referred to that paper for a discussion of the possible effects of dvuamical fiction ou the simulations.,"1998), we assume that the proto-Galactic fragments have equal merger probabilities; the reader is referred to that paper for a discussion of the possible effects of dynamical friction on the simulations."
+ Perhaps the most noteworthy feature of the simulations prescuted here is the diversity of the end-products: the sinulated globular cluster metallicity distributions show a wide range in appearance. ranging from uninodal distributions to more complex. ones having multiple distinct peaks.," Perhaps the most noteworthy feature of the simulations presented here is the diversity of the end-products: the simulated globular cluster metallicity distributions show a wide range in appearance, ranging from unimodal distributions to more complex ones having multiple distinct peaks."
+ This diversity is not unexpected iu a stochastic process such as galaxy formation. and differs from the predominantly bimodal globular cluster uectallhiitv distributions found previously for elaut elliptical galaxies (Cotté et al.," This diversity is not unexpected in a stochastic process such as galaxy formation, and differs from the predominantly bimodal globular cluster metallicity distributions found previously for giant elliptical galaxies (Côtté et al."
+ 1998) for two simple reasons., 1998) for two simple reasons.
+ First. bv virtue of the elobular cluster iietallicity-inuinositv relation and the modest hunünositv of the Galactic bulee. the mean metallicity of the imoetalxicli clusters is not as widely separated from that of the metal-»»or coniponent.," First, by virtue of the globular cluster metallicity-luminosity relation and the modest luminosity of the Galactic bulge, the mean metallicity of the metal-rich clusters is not as widely separated from that of the metal-poor component."
+ Secoud. the hieh huuinosities of eiaut elliptical galaxies permit the accretion of correspondiunglv nore luminous proto-galactic fragnieuts or galaxies. ucaning that the exponential cutoff in the hunuinositv distribution given by equation (5) iuiposes a sharp cutoff ou the metal-rich side of the distribution of elobular clusters arising iu proto-galactic fragments;," Second, the high luminosities of giant elliptical galaxies permit the accretion of correspondingly more luminous proto-galactic fragments or galaxies, meaning that the exponential cutoff in the luminosity distribution given by equation (5) imposes a sharp cutoff on the metal-rich side of the distribution of globular clusters arising in proto-galactic fragments."
+" Such a cutoff. which serves to dilineate the elobular clusters of the dominant proto-galactic fragment from those of the other fragments. does not apply in the case of the Calactic spheroid since Ly.<<Ly-. as discussed in §3.3 and 3.5,"," Such a cutoff, which serves to dilineate the globular clusters of the dominant proto-galactic fragment from those of the other fragments, does not apply in the case of the Galactic spheroid since $L_V^s \lae L_V^*$, as discussed in 3.3 and 3.5."
+ The principal conchisious drawn from these simulations can be sumunarized as follows: (1) the bulee. as the dominant proto-Galactic building block. is observed to have the most metal-rich globular clusters system by virtue ofthe globular cluster metallicitv-Iuninositv relation: (2) the bulee contributes roughlytwice the number of globular clusters initially as do the combined halo progenitors: (3) the more centrally concentrated metal-rich elobular cluster xvsteui las been preferentially eroded by dynamical effects: (C1) the metal-poor cluster svstemn exhibits a much wider range in its observed properties although the sinulatious reveal that a~9 produces the closest match to the observed distribution. particularly the peak at [Fe/TII| ~1.6: and (5) the metal-poor clusters. being more spatially exteuded— than their imietalaxicli counterparts. have mudereoue less severe dynamical erosion.," The principal conclusions drawn from these simulations can be summarized as follows: (1) the bulge, as the dominant proto-Galactic building block, is observed to have the most metal-rich globular clusters system by virtue of the globular cluster metallicity-luminosity relation; (2) the bulge contributes roughly the number of globular clusters initially as do the combined halo progenitors; (3) the more centrally concentrated metal-rich globular cluster system has been preferentially eroded by dynamical effects; (4) the metal-poor cluster system exhibits a much wider range in its observed properties although the simulations reveal that $\alpha \sim -2$ produces the closest match to the observed distribution, particularly the peak at [Fe/H] $\sim -1.6$; and (5) the metal-poor clusters, being more spatially extended than their metal-rich counterparts, have undergone less severe dynamical erosion."
+ The upper panel of Figure | shows a single simulation of elobular cluster metallicity distribution. The actual distribution. base onu a sample of 133 clusters having measured metallicities (Πας 1996b). is indicated by the open circles.," The upper panel of Figure 4 shows a single simulation of globular cluster metallicity distribution, The actual distribution, based on a sample of 133 clusters having measured metallicities (Harris 1996b), is indicated by the open circles."
+ The two dotted curves show the initial distributions of bulge aux halo elobular clusters. while the dashed curves iudicate the sale distributions after including the effects dywnuaniica erosiou.," The two dotted curves show the initial distributions of bulge and halo globular clusters, while the dashed curves indicate the same distributions after including the effects dynamical erosion."
+" The spheroid in this simulation has Z3.=7.1«109L4-.. aud contziis 119 surviving elobulax clusters. descended frou, au original population of 129."," The spheroid in this simulation has $L_V^s = 7.4\times10^9L_{V,{\odot}}$ and contains 149 surviving globular clusters, descended from an original population of 429."
+ Iu this particular simulation. the halo was assembled from a total of 1309 proto-Galactic fragiments having a combined DIuninositv of =27<JUL.," In this particular simulation, the halo was assembled from a total of 1309 proto-Galactic fragments having a combined luminosity of $L_V^h = 2.7\times10^9L_{V,{\odot}}$ ."
+" The final specific frequency of theLi. Galactic spheroid is $,=1.6. which is identical to the observed value."," The final specific frequency of the Galactic spheroid is $S_n = 1.6$, which is identical to the observed value."
+ The solid curve iu the lower panel of Figure { shows he halo star metallicity distribution for the simulation described above., The solid curve in the lower panel of Figure 4 shows the halo star metallicity distribution for the simulation described above.
+ Open circles indicate the actual netallicity distribution of 372 kinematically selected. Lalo Ποια stars according to Ryan Norris (1991)., Open circles indicate the actual metallicity distribution of 372 kinematically selected halo field stars according to Ryan Norris (1991).
+ Both the simulated aud observed distributions show maxima in the range 1.7< |Fe/T| L1.5 aud broad wings extending o lower aud higher inetallicities., Both the simulated and observed distributions show maxima in the range $-1.7 \lae$ [Fe/H] $\lae -1.5$ and broad wings extending to lower and higher metallicities.
+ These extended tails are nore pronouced in the field star distribution than in the cluster imetallicity distribution., These extended tails are more pronouced in the field star distribution than in the cluster metallicity distribution.
+ The siguificauce of these differences has always remained somewhat questionable due to the finite size of the Galactic cluster svstem (/.6.. the discrepancy at the metal-poor aud moetalaich cuds cau removed by adding oulv four aud six clusters. respectively: Laird ct al.," The significance of these differences has always remained somewhat questionable due to the finite size of the Galactic cluster system $i.e.$, the discrepancy at the metal-poor and metal-rich ends can removed by adding only four and six clusters, respectively; Laird et al."
+ 1985)., 1988).
+ Nevertheless. such differences are often seen in the siuulatiouns described here and their origi can be understood as follows.," Nevertheless, such differences are often seen in the simulations described here and their origin can be understood as follows."
+ The extended metal-poor tail in he field star distribution is populated exclusively by stars formed in the faintest aud most imetaldeficicut xoto-Calactie fragments: 56. having mean metallicities of [Fe/If| ~—2.1 (see Figure 5).," The extended metal-poor tail in the field star distribution is populated exclusively by stars formed in the faintest and most metal-deficient proto-Galactic fragments: $i.e.$, having mean metallicities of [Fe/H] $\sim -2.1$ (see Figure 5)."
+ Such a tail is slightly ess evident in the elobular cluster distribution since ouly proto-Galactic fragients haviug Ly=16x101σον. contribute clusters to the spheroid. limiting he iean ietallicity of these clusters to ΙΕΠΠ ~ 1.9.," Such a tail is slightly less evident in the globular cluster distribution since only proto-Galactic fragments having $L_V \gae 1.6\times10^7L_{V,{\odot}}$ contribute clusters to the spheroid, limiting the mean metallicity of these clusters to [Fe/H] $\sim -1.9$ ."
+ More significauth. the comparatively larec number of faint proto-Galactic fragments incorporated iuto the ido cusures that the metal-poor tail of simulated halo uctallicity distribution is well populated. iu coustrast to iat of the elobular cluster distribution (which. uulike the Ποια. star distribution. is further depleted by dynamical fects).," More significantly, the comparatively large number of faint proto-Galactic fragments incorporated into the halo ensures that the metal-poor tail of simulated halo metallicity distribution is well populated, in constrast to that of the globular cluster distribution (which, unlike the field star distribution, is further depleted by dynamical effects)."
+ At the ametalrich eud. the sinall umuber of uuinous proto-Galactic fragments incorporated iuto the spheroid contribute significant numbers of both globular clusters ancl field stars: as discussed in 83.2. these stars will be systematically ~ 0.6 dex more metal-rich than 1c associated clusters aud. consequently. will produce a uetalarich tail which will not be secu iu the globular cluster netallicity distribution.," At the metal-rich end, the small number of luminous proto-Galactic fragments incorporated into the spheroid contribute significant numbers of both globular clusters and field stars; as discussed in 3.2, these stars will be systematically $\sim$ 0.6 dex more metal-rich than the associated clusters and, consequently, will produce a metal-rich tail which will not be seen in the globular cluster metallicity distribution."
+ For comparison. the dotted curve in the lower paucl of Figure [| shows the prediction of the Simple Model or a effective vield of logiyye=—1.6. anarbitrary value chosen by Ryan Norris (1991) to give the closest natch to the observed distribution.," For comparison, the dotted curve in the lower panel of Figure 4 shows the prediction of the Simple Model for an effective yield of ${\log_{10}{y_e}} = -1.6$, an value chosen by Ryan Norris (1991) to give the closest match to the observed distribution."
+ As noted bv both Laid et al. (, As noted by both Laird et al. (
+1988) aud Ryan Norris (1991). at high uetallicities the Simple Model shows poor agreecimeut with the actual distrubutiou.,"1988) and Ryan Norris (1991), at high metallicities the Simple Model shows poor agreement with the actual distrubution."
+ While the significauce. of lis discrepancy is uuclear due to possible coutauinatiou woametal-rich disk stars. we note that the simulations show significant nunibers of stars having [Fe/I]| >1 whose origin can be traced to the largest proto-Galactic raesnients incorporated in the spheroid.," While the significance of this discrepancy is unclear due to possible contamination by metal-rich disk stars, we note that the simulations show significant numbers of stars having [Fe/H] $> -1$ whose origin can be traced to the largest proto-Galactic fragments incorporated in the spheroid."
+ The fraction of such stars in the simulations. however. is somewhat larger hanthat seen in the distribution of Ryan Norris (1991).," The fraction of such stars in the simulations, however, is somewhat larger thanthat seen in the distribution of Ryan Norris (1991)."
+ A consistency check ou the radial distribution of the, A consistency check on the radial distribution of the
+as we would expect.,as we would expect.
+" However, the region outside of r=0.25 is well described by the linear solution, with |Aout|«Ais|."," However, the region outside of $r\approx 0.25$ is well described by the linear solution, with $|A_{out}| \ll |A_{in}|$."
+" In this region, &.<1, so it is reasonable to assume that the IWs are efficiently absorbed near the centre."," In this region, $\mathcal{R} \ll 1$, so it is reasonable to assume that the IWs are efficiently absorbed near the centre."
+" This picture is identical to that in 2D. The picture in 3D in the ry-plane is very similar to that in 2D, as can be seen in Fig."," This picture is identical to that in 2D. The picture in 3D in the $xy$ -plane is very similar to that in 2D, as can be seen in Fig."
+ 5 (to compare with Fig., \ref{256lam15fr1t45xy1} (to compare with Fig.
+ 6 from BO10)., 6 from BO10).
+" However, one noticeable difference is that the primary wave preferentially transfers its angular momentum at low latitudes, close to the equatorial plane."," However, one noticeable difference is that the primary wave preferentially transfers its angular momentum at low latitudes, close to the equatorial plane."
+" This can be seen in Fig. 6,,"," This can be seen in Fig. \ref{256lam15fr1t45},"
+" where we plot the angular frequency of the flow normalised to 2, once a critical layer has formed, in both the zy and zz planes."," where we plot the angular frequency of the flow normalised to $\Omega_{p}$ once a critical layer has formed, in both the $xy$ and $xz$ planes."
+" This is a consequence of the latitudinal form of Y2, whose magnitude peaks at 0=«/2, as is illustrated in Fig. 3.."," This is a consequence of the latitudinal form of $Y_{2}^{2}$, whose magnitude peaks at $\theta = \pi/2$, as is illustrated in Fig. \ref{256lam15fr01t35xz}."
+" The critical layer absorption is observed to continue as the wave forcing is ongoing, so this differential rotation is continually reinforced by the absorption of |=m2 IWs."," The critical layer absorption is observed to continue as the wave forcing is ongoing, so this differential rotation is continually reinforced by the absorption of $l=m=2$ IWs."
+" Since there are wave motions in the region of fluid interior of the critical layer, parts of these regions spin slightly faster than Q, (this is not seen in Fig."," Since there are wave motions in the region of fluid interior of the critical layer, parts of these regions spin slightly faster than $\Omega_{p}$ (this is not seen in Fig."
+ 6 due to the adopted colour scale)., \ref{256lam15fr1t45} due to the adopted colour scale).
+ This was also observed in the 2D simulations., This was also observed in the 2D simulations.
+" The main motivation for our work is to study Q’ for solar- stars, and in particular to connect this with the survival of close-in extrasolar planets."," The main motivation for our work is to study $Q^{\prime}$ for solar-type stars, and in particular to connect this with the survival of close-in extrasolar planets."
+" We have demonstrated that the fate of IGWs approaching the centre of a solar-type star, outlined in BO10, is unaffected by the extension to three dimensions."," We have demonstrated that the fate of IGWs approaching the centre of a solar-type star, outlined in BO10, is unaffected by the extension to three dimensions."
+ A critical layer is formed by the deposition of angular momentum through wave breaking once the waves cause isentropic overturning near the centre., A critical layer is formed by the deposition of angular momentum through wave breaking once the waves cause isentropic overturning near the centre.
+" For lower amplitudes, the waves reflect coherently and approximately perfectly from the centre of the star, and may form global modes."," For lower amplitudes, the waves reflect coherently and approximately perfectly from the centre of the star, and may form global modes."
+" In that case, the dissipation is only efficient when the system enters a resonance with a global mode."," In that case, the dissipation is only efficient when the system enters a resonance with a global mode."
+" The corresponding time-averaged dissipation rate is weak, because the system spends most of its time out of resonance (T98; GD98), resulting in negligible tidal evolution of the planetary companion."," The corresponding time-averaged dissipation rate is weak, because the system spends most of its time out of resonance (T98; GD98), resulting in negligible tidal evolution of the planetary companion."
+" Note, however, that this neglects the possibility that passage through resonance will cause wave breaking, which should be studied in future work."," Note, however, that this neglects the possibility that passage through resonance will cause wave breaking, which should be studied in future work."
+" If a critical layer forms, wave absorption is efficient, and global modes (of any frequency very similar to the orbital frequency) are prevented from being set up in the RZ."," If a critical layer forms, wave absorption is efficient, and global modes (of any frequency very similar to the orbital frequency) are prevented from being set up in the RZ."
+ The tidal torque can then be computed from assuming that the [Ws are entirely absorbed., The tidal torque can then be computed from assuming that the IWs are entirely absorbed.
+ In this case a calculation, In this case a calculation
+"formation of our Galaxy. ideutifv when cluster formation occurred with respect to the era of reionization iu the Universe and the cosuuc star formation peak. constrain the formation history of the Galaxy by dating the halo and bulge. as well as constrain O3; (Crattonctal.2003). when coupled with measurements of £7, bv WALAP (Sperecletal.2003) and the IST Nev Project (Freedinanetal.2001).","formation of our Galaxy, identify when cluster formation occurred with respect to the era of reionization in the Universe and the cosmic star formation peak, constrain the formation history of the Galaxy by dating the halo and bulge, as well as constrain $\Omega_M$ \citep{gratton03} when coupled with measurements of $H_{\rm{o}}$ by WMAP \citep{spergel03} and the HST Key Project \citep{freedman01}."
+. Most. of the Galactic globular clusters have ages between 9-114 Cyr (e.g.Marin-Franchetal.2009:Dotterct 2010).. with most absolute ages cohnunouly determined by fitting theoretical isochrones to the main sequence turnoff.," Most of the Galactic globular clusters have ages between 9-14 Gyr \citep[e.g.][]{marinfranch09,dotter10}, with most absolute ages commonly determined by fitting theoretical isochrones to the main sequence turnoff."
+ This. however. depeuds strongly on the assmned distance modulus to the cluster. which is the most significant contributor to the age uneertainty (Bolte&Toean1995).," This, however, depends strongly on the assumed distance modulus to the cluster, which is the most significant contributor to the age uncertainty \citep{bolte95}."
+. Althoueh less sensitive to distance. even the more receuth eniploved age dating technique via the white dwarf cooling sequence (Hausenetal.2001.2007). requires a precise distance modulus.," Although less sensitive to distance, even the more recently employed age dating technique via the white dwarf cooling sequence \citep{hansen04,hansen07} requires a precise distance modulus."
+ The Calactic globular cluster [7 Tucanac (17 Tuc. NGC 101) is considered the metalrich proto-type for elobular cluster work (|Fe/I]]= 00.026..," The Galactic globular cluster 47 Tucanae (47 Tuc, NGC 104) is considered the metal-rich proto-type for globular cluster work \citep[$\rm{[Fe/H]}= 0.026,."
+—2009). As the closest aud least reddened metal-rich globular cluster. £7 Tuc serves as au iuportant science. as well as a calibration target. for the IIubble Space Telescope (IST).," As the closest and least reddened metal-rich globular cluster, 47 Tuc serves as an important science, as well as a calibration target, for the Hubble Space Telescope (HST)."
+ Tt is used as a metalvich anchor when comparing observations of resolved stars In nearby ealaxics to stellar evolutionary models (Cassis?&Salaris1997:Zoccalietal.1999:Ferraro1999:DiCoccoetal.2010). and in examining the star formation history of galaxies in the Local Croup (e.g.Monellietal. 2010).," It is used as a metal-rich anchor when comparing observations of resolved stars in nearby galaxies to stellar evolutionary models \citep[][]{cassisi97,zoccali99,ferraro99,dicecco10} and in examining the star formation history of galaxies in the Local Group \citep[e.g.][]{monelli10}."
+.. As another example. the Wide Field Camera 3 (WFEC3) Galactic Bulee Treasury. program (CO-L1661) has targeted. £7 Tuc in their resolved stellar population study (Brownctal.2009).," As another example, the Wide Field Camera 3 (WFC3) Galactic Bulge Treasury program (GO-11664) has targeted 47 Tuc in their resolved stellar population study \citep{brown09}."
+. With these observations. ranging from the UV. to near-IR. 17 Tuc will serve as an curpirical population template. helping," With these observations, ranging from the UV to near-IR, 47 Tuc will serve as an empirical population template, helping"
+"Recently. the PDF of the Tsallis ensemble was linked to (he analysis of fully. developed (turbulence providing a relation between nonextensivity ancl intermittency, which manilests the multifraetalitv in the distribution of eddies (Arvimitsn&Arimitsn2000a.2001).","Recently, the PDF of the Tsallis ensemble was linked to the analysis of fully developed turbulence providing a relation between nonextensivity and intermittency, which manifests the multifractality in the distribution of eddies \citep{Arimitsu00a,Arimitsu01}."
+. It was shown that (he value of the index 4 of nonextensive statistics is related. to the extremes of the multifractal spectrum. (Lyra.&Tsallis1998). and that the multifractal spectrum corresponding to the Tsallis statistics is determined self-consistently (Arimitsu&Arimitsu 2000b)., It was shown that the value of the index $q$ of nonextensive statistics is related to the extremes of the multifractal spectrum \citep{Lyra98} and that the multifractal spectrum corresponding to the Tsallis statistics is determined self-consistently \citep{Arimitsu00b}.
+. Moreover. the context of generalized (hermo-stalislics provides analvtical formulas for PDFs of distance dependent velocity differences. implying the cascade like structure of the turbulent dvnanmies ancl an interpretation of the parameter &=1/(q—1) (Beck2000).," Moreover, the context of generalized thermo-statistics provides analytical formulas for PDFs of distance dependent velocity differences, implying the cascade like structure of the turbulent dynamics and an interpretation of the parameter $\kappa = 1/(q-1)$ \citep{Beck00}."
+. Links between (he nonextensive parameter q and (he corresponding parameters of the log-normal. multifractal and random-? models were derived (Shivamogei&Beck2003).," Links between the nonextensive parameter $q$ and the corresponding parameters of the log-normal, multifractal and $\beta$ models were derived \citep{Shivamoggi03}."
+. We relate. in the following. nonlocalitv in turbulent [lows to nonextensive svstems andl demonstrate in (he context of entropy. generalization the consistency. of (he theoretical bi- distribution (Leubner2004a) wilh observed. scale dependent PDFs of characteristic Κον variables in the intermittent. turbulent interplanetary medium. where & appears as the only tuning parameter between the scales.," We relate, in the following, nonlocality in turbulent flows to nonextensive systems and demonstrate in the context of entropy generalization the consistency of the theoretical bi-kappa distribution \citep{Leubner04a} with observed, scale dependent PDFs of characteristic key variables in the intermittent, turbulent interplanetary medium, where $\kappa$ appears as the only tuning parameter between the scales."
+ The nonextensive entropy generalization for the classical thermo-statistics proposed bv Tsallis(1988). takes the form where p; is the probability of the i”i microstate. j is Boltzmann's constant and q is a parameter quantilving the degree of nonextensivitv of the system. commonly referred (o as the entropic index.," The nonextensive entropy generalization for the classical thermo-statistics proposed by \citet{Tsallis88} takes the form where $p_{i}$ is the probability of the $i^{th}$ microstate, $k_{B}$ is Boltzmann's constant and $q$ is a parameter quantifying the degree of nonextensivity of the system, commonly referred to as the entropic index."
+ A crucial property of this entropy is the pseuclo-acditivity such that for given subsystems d and 2 in the sense of [actorizability of the microstate probabilities., A crucial property of this entropy is the pseudo-additivity such that for given subsystems $A$ and $B$ in the sense of factorizability of the microstate probabilities.
+ llence. nonlocality or long-range interactions are introduced by (he multüiplicative term accounting lor correlations between the subsystems.," Hence, nonlocality or long-range interactions are introduced by the multiplicative term accounting for correlations between the subsystems."
+ In order to link the Tsallis q—statistics to the family of &—distributions applied in space and astroplivsical plasma modeling we perform the (transformation, In order to link the Tsallis $q-$ statistics to the family of $\kappa-$ distributions applied in space and astrophysical plasma modeling we perform the transformation
+Γονοί].,beyond.
+" The restriction to medium resolution spectroscopy in the blue spectral range provides significant advantages,", The restriction to medium resolution spectroscopy in the blue spectral range provides significant advantages.
+ Most importantly. contamination flrom sky ancl region emission is bv far less critical than in the red. which is needed as an additional spectral range. for instance. lor the WLR method. which requires the measurements of Ho profiles with at least ~QA resolution.," Most importantly, contamination from sky and region emission is by far less critical than in the red, which is needed as an additional spectral range, for instance, for the WLR method, which requires the measurements of $\alpha$ profiles with at least $\sim2\,$ resolution."
+ Moreover. (he amount of observing time for an accurate distance determination is significantly reduced. if only spectra in (he blue are required.," Moreover, the amount of observing time for an accurate distance determination is significantly reduced, if only spectra in the blue are required."
+ The accurate calibration of the FGLR (and the WLR) will become the crucial element of future work. before the method can be applied seriously for extragalactic distance determinations.," The accurate calibration of the FGLR (and the WLR) will become the crucial element of future work, before the method can be applied seriously for extragalactic distance determinations."
+ Local Group galaxies with well determined distances provide the ideal laboratory for this purpose., Local Group galaxies with well determined distances provide the ideal laboratory for this purpose.
+ Multi-object spectrographs with rather high spectral resolution attached {ο 8 to LOm-class telescopes such as FLAMES (VLT) or DEIMOS (Ixeck 2) will allow high quality spectra of large candidate samples in each galaxy. with a rather modest amount of observing (ime., Multi-object spectrographs with rather high spectral resolution attached to 8 to 10m-class telescopes such as FLAMES (VLT) or DEIMOS (Keck 2) will allow high quality spectra of large candidate samples in each galaxy with a rather modest amount of observing time.
+ such a svstematic study of hundreds of blue supergiants in Local Group galaxies will 10b Only provide an accurate calibration of the FGLR., Such a systematic study of hundreds of blue supergiants in Local Group galaxies will not only provide an accurate calibration of the FGLR.
+ It will also enable us to investigate important aspects of stellar evolution. which are related.," It will also enable us to investigate important aspects of stellar evolution, which are related."
+ The most crucial ones concern the role of metallicity and stellar rotation., The most crucial ones concern the role of metallicity and stellar rotation.
+ Investigating stellar evolutionary (racks at different nuetallicitv and with dillerent initial rotation at (he main sequence (AlevnetMevnetetal. 1994)) we find small. but noticeable effects on the theoretical ΕΤΗ.," Investigating stellar evolutionary tracks at different metallicity and with different initial rotation at the main sequence \citealt{meynet00,meynet94}) ) we find small, but noticeable effects on the theoretical FGLR."
+ Mass-loss in evolutionary stages prior to the blue supergiant phase depends on metallicity 2000)) and has a small influence on the FOLR., Mass-loss in evolutionary stages prior to the blue supergiant phase depends on metallicity \citealt{kud00}) ) and has a small influence on the FGLR.
+ Rotation affects the strength of and the internal mixing processes and introduces a mocification of the mass-Iuminositv relationship in the blue supereiant stage., Rotation affects the strength of mass-loss and the internal mixing processes and introduces a modification of the mass-luminosity relationship in the blue supergiant stage.
+ The effect of rotation becomes larger at higher unminosities., The effect of rotation becomes larger at higher luminosities.
+ This might be the reason for the increase of the residual scatter at luminosities above M4;=—8 mentioned above., This might be the reason for the increase of the residual scatter at luminosities above $M_{bol}=-8$ mentioned above.
+ Another very important issue is the fraction among the sample of observed blue supergiants evolving backwards to the blue aller a previous phase as red supergiants., Another very important issue is the fraction among the sample of observed blue supergiants evolving backwards to the blue after a previous phase as red supergiants.
+ Those objects are expected to have lost a significant [raction of their mass as red supergiants and might form an additional sequence below (he observed relationship., Those objects are expected to have lost a significant fraction of their mass as red supergiants and might form an additional sequence below the observed relationship.
+ Evolutionary calculations indicate that the relative nunber of those objects might depend crucially on metallicity and rotation., Evolutionary calculations indicate that the relative number of those objects might depend crucially on metallicity and rotation.
+ Systematic studies of Local Group galaxies at different metallicity. as proposed above. will allow to investigate this problem.," Systematic studies of Local Group galaxies at different metallicity, as proposed above, will allow to investigate this problem."
+ In general. (he observational detection of the Geht relationship between flux-weishted," In general, the observational detection of the tight relationship between flux-weighted"
+lighest welocities (terminal side).,highest velocities (terminal side).
+ These. along with nünor axis (bv) positiou-velocity diagrams. as well as channel maps of models aud data must all be examined to fine tune the above parameters iu the models. while at the same time ensuriug that the vertical profiles and zerotl-noment maps are well matched.," These, along with minor axis (bv) position-velocity diagrams, as well as channel maps of models and data must all be examined to fine tune the above parameters in the models, while at the same time ensuring that the vertical profiles and zeroth-moment maps are well matched."
+ Although all plots are considered throughout the process. the by diagrams aud chaunel maps are the most heavily relied upon after the initial modeling stages.," Although all plots are considered throughout the process, the bv diagrams and channel maps are the most heavily relied upon after the initial modeling stages."
+ As will be deinonstrated. the kev to this method is to start with the simplest possible model aud then add features such as warps aud flares - one at a time - to fully uuderstaud the individual contribution of each additional degree of complexity.," As will be demonstrated, the key to this method is to start with the simplest possible model and then add features such as warps and flares - one at a time - to fully understand the individual contribution of each additional degree of complexity."
+ This allows for subtle distinctious between features. which otherwise may be mismnterpreted.," This allows for subtle distinctions between features, which otherwise may be misinterpreted."
+ Iu the end. au acceptable model is oue that fits of the plots well.," In the end, an acceptable model is one that fits of the plots well."
+ The quoted errors indicate visually estimated uncertainties rather than formal lo error bars., The quoted errors indicate visually estimated uncertainties rather than formal $\sigma$ error bars.
+ It should be noted that new. seni-autoiated tilted ving fitting software (TiRiFiC) (?7).. which allows for a chi-square fit of models to observed data cubes is now available.," It should be noted that new, semi-automated tilted ring fitting software ) \citep{2007A&A...468..731J}, which allows for a chi-square fit of models to observed data cubes is now available."
+ Advanced models for NOC 1211 were examined in bothTiRiFiC aud aud cach produced comparable results., Advanced models for NGC 4244 were examined in both and and each produced comparable results.
+ Caven the advanced stage of modeling already completed viagalmod. as well as the localized. features in NGC 1211. was not usec extensively in this case.," Given the advanced stage of modeling already completed via, as well as the localized features in NGC 4244, was not used extensively in this case."
+ However. the adveut of TÀiRiFiCs automated processes will help to expedite the fitting of the remaining ealaxies iu the TALOGAS sample.," However, the advent of 's automated processes will help to expedite the fitting of the remaining galaxies in the HALOGAS sample."
+ The by diagrams shown in Figure 6 and the chauncl maps in Fieure 7 illustrate a number of the features that were noticeably scusitive to the model parameters we endeavor to match., The bv diagrams shown in Figure \ref{fig6} and the channel maps in Figure \ref{fig7} illustrate a number of the features that were noticeably sensitive to the model parameters we endeavor to match.
+" There is a notable ""T shape in the panels closest to the galactic ceuter iu the bv plots in Figure G.. which morphs into a ""V shape in the outer panels (schematically indicated in certain pancls of Fieure 6))."," There is a notable “T"" shape in the panels closest to the galactic center in the bv plots in Figure \ref{fig6}, which morphs into a “V"" shape in the outer panels (schematically indicated in certain panels of Figure \ref{fig6}) )."
+ Additionally. the contours ou tle svstemic side are flattened.," Additionally, the contours on the systemic side are flattened."
+ There is also a substantial lopsideduess vest seen ato 5.155 aud 6.9 iu Figure 6 around the uidplaue. primarily due to the component of the warp xrpendieular to the line of sight. but also partiallydue ο asvintnctrics.," There is also a substantial lopsidedness best seen at $-$ 5.1' and $-$ 6.9' in Figure \ref{fig6} around the midplane, primarily due to the component of the warp perpendicular to the line of sight, but also partiallydue to asymmetries."
+ The channel maps displaved in Figure 7 reiterate he need for a warp perpendicular to the line of sieht., The channel maps displayed in Figure \ref{fig7} reiterate the need for a warp perpendicular to the line of sight.
+ Prominent features that were sensitive to modcl »uwnneters considered during the modehug process include the spacing and anele of the contours at the ips of each plot closest to the ceuter of the galaxy., Prominent features that were sensitive to model parameters considered during the modeling process include the spacing and angle of the contours at the tips of each plot closest to the center of the galaxy.
+ Also considered is the approximately 15 deerce slaut along the outermost edges at more extreme velocities aud at mostly iegative nuünor axis offsets in the approaching half arc »ositive offsets in the recediug half., Also considered is the approximately 45 degree slant along the outermost edges at more extreme velocities and at mostly negative minor axis offsets in the approaching half and positive offsets in the receding half.
+" Iu both Figures G and τν, the visual assessment of he quality of the fit between the models aud data is indicated in each panel via svinbols explained in the caption of Figure 6.."," In both Figures \ref{fig6} and \ref{fig7}, the visual assessment of the quality of the fit between the models and data is indicated in each panel via symbols explained in the caption of Figure \ref{fig6}."
+ The assiguiment of cach sviubo reavily relies on the criteria described above., The assignment of each symbol heavily relies on the criteria described above.
+ Some quantities are kept constaut throughout all of he moclels for the entire process., Some quantities are kept constant throughout all of the models for the entire process.
+" These are the svstemiic velocity (211 kins 13. kinematic ceuter (12h17129.90s, 37dIS129.00s) and the run of the position angle derive roni the component of the warp perpendicular to the ine of sight. cohunu density profile. rotation curve (both shown in Figure 8)) aud velocity dispersion (12 kin E decreasing to 10 knis b ata radius of. 7.7 or 10 kpe)."," These are the systemic velocity (244 km $^{-1}$ ), kinematic center (12h17m29.90s, 37d48m29.00s) and the run of the position angle derived from the component of the warp perpendicular to the line of sight, column density profile, rotation curve (both shown in Figure \ref{fig8}) ) and velocity dispersion (12 km $^{-1}$ decreasing to 10 km $^{-1}$ at a radius of 7.7' or 10 kpc)."
+ Several possible models are considered., Several possible models are considered.
+ Some are Claninated quickly upon inspection| of the vertical profile or zeroth-oment map. while others require substantially iuore effort to discern between them.," Some are eliminated quickly upon inspection of the vertical profile or zeroth-moment map, while others require substantially more effort to discern between them."
+ Vertical distributions are expoucutials im all models., Vertical distributions are exponentials in all models.
+ Expoucutial scale heights vary among models aud are listed in Table 3..., Exponential scale heights vary among models and are listed in Table \ref{tbl-3}. .
+ All of the models. uuless otherwise," All of the models, unless otherwise"
+Tn order to check this idea. we performed a single N-body simulation of this capture.,"In order to check this idea, we performed a single N-body simulation of this capture."
+ The code used was the tree-code from Darues Tut (1986... 19893).," The code used was the tree-code from Barnes Hut \cite{Barnes1986}, \cite{Barnes1989}) )."
+ Table 1 sunuuuizes the initial conditious. ie. the scale-leneth and total mass of cach component.," Table 1 summarizes the initial conditions, i.e. the scale-length and total mass of each component."
+ All components were truncated at a radius of 32 kpc., All components were truncated at a radius of 32 kpc.
+ The main galaxy (simulating the Iate-tvpe spiral NCC 5907) was composed of a spherical non-rotatiug bulee.o represented by a Plunuuer component. an n»= Tooure stellar disk. thickened by a sech z-distribution with a coustant scale height of zy=500 pe. and a spherical Plummer halo.," The main galaxy (simulating the late-type spiral NGC 5907) was composed of a spherical non-rotating bulge, represented by a Plummer component, an $n=1$ Toomre stellar disk, thickened by a $^2$ z-distribution with a constant scale height of $z_0 = 500$ pc, and a spherical Plummer halo."
+ The program computes the combined potential of all componcuts for the main galaxy. and solves the Jeans equations to derive the initial rotational velocity aud dispersion.," The program computes the combined potential of all components for the main galaxy, and solves the Jeans equations to derive the initial rotational velocity and dispersion."
+ The stellar disk was launched with a Tooure piruueter Q slightly decreasing with radius. from 1.5 at) the center to 1 at the eud of the disk.," The stellar disk was launched with a Toomre parameter $Q$ slightly decreasing with radius, from 1.5 at the center to 1 at the end of the disk."
+ Initially the disk is mnildy unstable against spiral aud bar formation., Initially the disk is mildy unstable against spiral and bar formation.
+ The resulting rotation curve is compatible with the observed one (Sancisi van Albada 1987)). as shown iu fimeure 2.," The resulting rotation curve is compatible with the observed one (Sancisi van Albada \cite{Sancisi}) ), as shown in figure 2."
+ The companion is represented by a spherical παπι! component. without dark halo.," The companion is represented by a spherical Plummer component, without dark halo."
+ Its mass is one teuth of that of 55907., Its mass is one tenth of that of 5907.
+" Tt is launched at 32 Ipc radius roni 55907, with a purely tanecutial velocity of 200 au Ll. du a bound. almost circular orbit."," It is launched at 32 kpc radius from 5907, with a purely tangential velocity of 200 km $^{-1}$, in a bound, almost circular orbit."
+ The initial configuration is displaved iu fie., The initial configuration is displayed in fig.
+ 3., 3.
+ Fig., Fig.
+ L| shows soe steps of the calculation: 300 Myr after the beeimuine of je Sinmlation. the elliptical galaxy has triggered a spiral structure In 55907 iud starts to bo slowed down by dyaiunical friction.," 4 shows some steps of the calculation: 300 Myr after the beginning of the simulation, the elliptical galaxy has triggered a spiral structure in 5907 and starts to be slowed down by dynamical friction."
+ Mereius is well advanced 1 Cyr after 16 beginning. and we stop the simulation after 1.9 Civi.," Merging is well advanced 1 Gyr after the beginning, and we stop the simulation after 1.9 Gyr."
+ At this time. the spiral galaxy has restored its flat disk structure. although it is warped and thicker due to the dynamical heating of its stars in the process.," At this time, the spiral galaxy has restored its flat disk structure, although it is warped and thicker due to the dynamical heating of its stars in the process."
+ This disk is tilted with respect to the original plane due to the conservation of total augular momentum., This disk is tilted with respect to the original plane due to the conservation of total angular momentum.
+ The elliptical COnpauidou has become a fat. hollow and distorted oblate ellipsoid. ceutered on the spiral galaxy aud with a rather siuuilar plane of svinmetry.," The elliptical companion has become a fat, hollow and distorted oblate ellipsoid centered on the spiral galaxy and with a rather similar plane of symmetry."
+ This final stage is displaved on fig., This final stage is displayed on fig.
+ 5., 5.
+ Fie., Fig.
+ 6 shows a cut through the “visible” particles of both galaxies. perpendicular to the new major axis of the spiral.," 6 shows a cut through the ""visible"" particles of both galaxies, perpendicular to the new major axis of the spiral."
+ The elliptical companion dominates the visible mass at heights larger than [| kpc from the plane. aud there the colors should be close to those of the elliptical ealaxy as observed.," The elliptical companion dominates the visible mass at heights larger than 4 kpc from the plane, and there the colors should be close to those of the elliptical galaxy as observed."
+ We cuphasize that the companion chosen for the experiment is 5-10 times more massive that what is needed to explain the observations: this choice was necessary to remedy the simall dvuauiucal range of N-body simulations. but is still indicative of the ανασα] phenomenon.," We emphasize that the companion chosen for the experiment is 5-10 times more massive that what is needed to explain the observations; this choice was necessary to remedy the small dynamical range of N-body simulations, but is still indicative of the dynamical phenomenon."
+ Even assuming a coustaut M/L for the stars of the elliptical. fe.," Even assuming a constant M/L for the stars of the elliptical, fig."
+ 6 shows that the helt docs not follow the dark matter profile. although the dark halo density is donuünant iu this region.," 6 shows that the light does not follow the dark matter profile, although the dark halo density is dominant in this region."
+ Iu auw case. there are maux choices for the r- aud z-distributioun of the dark matter. even a rotation curve. and this experiment shows that the light can be a poor tracer of the dark matter.," In any case, there are many choices for the r- and z-distribution of the dark matter, given a rotation curve, and this experiment shows that the light can be a poor tracer of the dark matter."
+ Tn this paper. we reported on new observations made with ιο CFII telescope which eoufixin cutirely the conclusious reached iu Paper E the luminous halo around 5907 is redder than the galaxy. and has the colors of a inetalrich old stellar population.," In this paper, we reported on new observations made with the CFH telescope which confirm entirely the conclusions reached in Paper I: the luminous halo around 5907 is redder than the galaxy and has the colors of a metal--rich old stellar population."
+ ITowever discrepancies betweeu the preseut observations aud others in the literature have fo be understood. which require further observations.," However discrepancies between the present observations and others in the literature have to be understood, which require further observations."
+" We have shown by a Nbody simulation that this red halo could result from cauuibalisii of a lowmass elliptical (a few 10? AL, after a slow encouuter with S907.", We have shown by a N–body simulation that this red halo could result from cannibalism of a low–mass elliptical (a few $^9$ $_\odot$ ) after a slow encounter with 5907.
+ Tow frequent is this phenomenon?, How frequent is this phenomenon?
+ Nearby spiral galaxies like M 31 often have elliptical coumpanious which iight eventually meree with them., Nearby spiral galaxies like M 31 often have elliptical companions which might eventually merge with them.
+ Our Calaxy iuieht have canmibalized a satellite which would have produced the thick disk as the result of stellar dvuamical heating. the process evidenced by our simulation (Robin ct al. 1996)).," Our Galaxy might have cannibalized a satellite which would have produced the thick disk as the result of stellar dynamical heating, the process evidenced by our simulation (Robin et al. \cite{Robin}) )."
+ To ect statistical information on this capture, To get statistical information on this capture
+facility of the National Science Foundation operated uncer Cooperative agreement by Associated Universities. Inc. This research has mace use of data obtained from the SuperCOSMOS Sky Surveys (SSS). prepared ancl hosted by the Wide Field Astronomy Unit. Institute for Astrononiv. University of Eclinbureh. which is funded by the Ulx Science and Technology Facilities Council.,"facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. This research has made use of data obtained from the SuperCOSMOS Sky Surveys (SSS), prepared and hosted by the Wide Field Astronomy Unit, Institute for Astronomy, University of Edinburgh, which is funded by the UK Science and Technology Facilities Council."
+ Phe SSS Web Site isSR., The SSS Web Site is.
+ The research also mace use of data from the SDSS survey., The research also made use of data from the SDSS survey.
+ Funding for the SDSS and SDSS-LL has been provided by the Alfred PL Sloan Foundation. the Participating Institutions. the National Science Foundation. the U.S. Department of Energy. the National Aeronautics and Space Administration. the Japanese Monbukagakusho. the Max Planck Society. and the Higher Education Funding Council for England.," Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England."
+ Phe SDSS Web Site isweee., The SDSS Web Site is.
+sdss.org/. Given an integral souree-count of the form No=AS integration time / per pointing. and a telescope settle time νι there is an optimum observing strategy to detect. the maximum number of sources.," Given an integral source-count of the form $N = K S^{-\gamma}$, integration time $t$ per pointing, and a telescope settle time $t_s$, there is an optimum observing strategy to detect the maximum number of sources."
+ It is easy to show that in the radio regime the is optimal: wide and shallow easily beats deep and narrow., It is easy to show that in the radio regime the is optimal; wide and shallow easily beats deep and narrow.
+ In a given time T we can make 7Z/(4|(i) pointings., In a given time T we can make $T/(t+t_s)$ pointings.
+ Consider the case for long integrations in. which ἐς is negligible., Consider the case for long integrations in which $t_s$ is negligible.
+" Phen the number of pointings is T/t. and if the equivalent beam area is zl, then the area covered. is A=AY."," Then the number of pointings is T/t, and if the equivalent beam area is $A_o$ then the area covered is $A=A_oT/t$."
+" The tus density limit becomes S=S,/Vf. where 5, is the Dux density reached in unit time."," The flux density limit becomes $S=S_{o} /
+\sqrt{t}$, where $S_o$ is the flux density reached in unit time."
+ Phe total number of sources found is N=AWSτσVID or NxTU?Us ," The total number of sources found is $N = AK S^{-\gamma} = T/t (S_{o} /\sqrt{t})^{-\gamma}$, or $N
+\propto T t^{(\gamma/2 - 1)}$."
+Dus until 5>2. until the source count issfeeper than that even at highest radio (us densities. the number of sources is maximized with minimum integration times and maximum number of pointings.," Thus until $\gamma > 2$, until the source count is than that even at highest radio flux densities, the number of sources is maximized with minimum integration times and maximum number of pointings."
+ In fact at the Dux densities in question 1., In fact at the flux densities in question $\gamma \sim 1$.
+ Given then that short exposures are required. what is optimum exposure time. in that ἐς must now come into play?," Given then that short exposures are required, what is optimum exposure time, in that $t_s$ must now come into play?"
+ There must be an optimum as if we let f+ 0. ἐν in itself imposes a maximum number of pointings forzero integration time.," There must be an optimum as if we let $t
+\rightarrow 0$ , $t_s$ in itself imposes a maximum number of pointings for integration time."
+ So let /= ru., So let $t = xt_s$ .
+ The limiting Εις density reached. is S=SfE the area covered in given observing time T (in which Z7/(E|]a) pointings are possible) is AZ76.1|ur).," The limiting flux density reached is $S
+= S_o/\sqrt{t}$; the area covered in given observing time $T$ (in which $T/t(1+x)$ pointings are possible) is $A_o T / t_s(1+x)$."
+ The total number of sources detected will be maximized at For Ξ1 the result. is particularly simple: russ=1. integration time and settle time should. be equal.," The total number of sources detected will be maximized at For $\gamma = 1$ the result is particularly simple: $x_{\rm max} = 1$, integration time and settle time should be equal."
+ Fig AL shows the function of equation T., Fig \ref{fig2_1} shows the function of equation 1.
+ Xt iezcris the function falls very slowly indeed., At $x \ge x_{\rm max}$ the function falls very slowly indeed.
+ This insensitivity tor suggests that a further factor could be considered - how much data do we not need?, This insensitivity to $x$ suggests that a further factor could be considered - how much data do we not need?
+" For ur=3 we achieve only half the pointings as [or .""=1. and vet the loss of sources is just 15 per cent."," For $x = 3$ we achieve only half the pointings as for $x=1$, and yet the loss of sources is just 15 per cent."
+ The NVSS used a 30s evele: 23s integration with 7s settle. w=3.3. representing a eood compromise.," The NVSS used a 30s cycle; 23s integration with 7s settle, $x = 3.3$, representing a good compromise."
+ The final factor to consider is whether to overlap the beams to make the survey complete to a limiting (lux density or to have them entirely independent in a coarse grid., The final factor to consider is whether to overlap the beams to make the survey complete to a limiting flux density or to have them entirely independent in a coarse grid.
+ The argument of the opening paragraph pertains - the answer is to use completely independent. beams to vield. maximum area at minimum: sensitivity., The argument of the opening paragraph pertains - the answer is to use completely independent beams to yield maximum area at minimum sensitivity.
+ Lhe issue can be quantified. quoting the analysis communicated to us by Jim Concon.," The issue can be quantified, quoting the analysis communicated to us by Jim Condon."
+ The NVSS spacing is overlapped just enough. to. provide nearly uniform sensitivity to ensure down to a [ixed limit (ο... 2.5 midv/beam) over areas 72 beam solid angle.," The NVSS spacing is overlapped just enough to provide nearly uniform sensitivity to ensure down to a fixed limit (e.g., 2.5 mJy/beam) over areas $>>$ beam solid angle."
+ To maximize the number of detections in a sample without achieving completeness. the optimum spacing is anvthing laree enough that successive snapshots have no overlap (1.6. >>2 FWHIAL).," To maximize the number of detections in a sample without achieving completeness, the optimum spacing is anything large enough that successive snapshots have no overlap (i.e. $> 2$ FWHM)."
+ Then. instead of a survey with uniform sensitivity. the vield is a survey covering a small area with high sensitivity and larger areas with lower sensitivity.," Then, instead of a survey with uniform sensitivity, the yield is a survey covering a small area with high sensitivity and larger areas with lower sensitivity."
+ Phe NVSS spacing gives the same sensitivity as the on-axis sensitivity. of a single snapshot over an elfective area of 1/2 the beam solid. angle. or z87/(SIn2) for a Gaussian beam with FWLIIM 8 (Condon et al.," The NVSS spacing gives the same sensitivity as the on-axis sensitivity of a single snapshot over an effective area of 1/2 the beam solid angle, or $\pi \theta^2 / (8
+{\rm ln} 2)$ for a Gaussian beam with FWHM $\theta$ (Condon et al."
+ 1998)., 1998).
+ A single isolated snapshot has sensitivity directly proportional to the primary power pattern., A single isolated snapshot has sensitivity directly proportional to the primary power pattern.
+ Since the source counts are roughly power-law with a dillerential slope of 2.0 (integral slope of 1.0). the number of detectable sources in the isolated snapshot will be proportional to the whole beam solid angle =x87/(41n2).," Since the source counts are roughly power-law with a differential slope of $2.0$ (integral slope of $1.0$ ), the number of detectable sources in the isolated snapshot will be proportional to the whole beam solid angle $= \pi \theta^2 / (4
+{\rm ln} 2)$."
+ Thus the viele per unit time with non-overlapping snapshots is that with snapshots overlapped as for the NVSS ericleling., Thus the yield per unit time with non-overlapping snapshots is that with snapshots overlapped as for the NVSS gridding.
+ Accordingly we adopted the NVSS 30s evcle. Ts settle and23s integration. andcoarse gridding. (," Accordingly we adopted the NVSS 30s cycle, 7s settle and23s integration, andcoarse gridding. ("
+A shorter settle time for the VLA would have been an advantage to us.),A shorter settle time for the VLA would have been an advantage to us.)
+Whatever the nature of the lens. if it is not nearby. but is instead close to the lensecl source star. [inite-source-size elfects are more likely to be detectable.,"Whatever the nature of the lens, if it is not nearby, but is instead close to the lensed source star, finite-source-size effects are more likely to be detectable."
+ Finite-source-size can be used to measure (he Einstein angle. θ).. of the lens. (hus providing a relationship between the lens mass and its distance. Dj trom us.," Finite-source-size can be used to measure the Einstein angle, $\theta_E$, of the lens, thus providing a relationship between the lens mass and its distance, $D_L$ from us."
+ Together with the value of τι derived [rom a fit to the light curve. θε can be used to compute the value of (he angular speed w=v/Dj. We can then test models in which both ¢ and {2 are relatively small. or relatively large.," Together with the value of $\tau_E$ derived from a fit to the light curve, $\theta_E$ can be used to compute the value of the angular speed $\omega=v/D_L.$ We can then test models in which both $v$ and $D_L$ are relatively small, or relatively large."
+ Finite-source-size effects also allow us to explore features of the source stars surface., Finite-source-size effects also allow us to explore features of the source star's surface.
+ In 2 we establish that an event with short Einstein-diameter crossing time can be caused only by a limited set of lenses., In 2 we establish that an event with short Einstein-diameter crossing time can be caused only by a limited set of lenses.
+ Section 3 is an overview of the (vpes of measurements that can be used to identilv the correct. physical model for the lens producing a event., Section 3 is an overview of the types of measurements that can be used to identify the correct physical model for the lens producing a short-duration event.
+ We focus on effects other than caustic crossings. which were studied for wide-orbit planets in Han 2005.," We focus on effects other than caustic crossings, which were studied for wide-orbit planets in Han 2005."
+ We show that in many cases the Jens mass anc its distance [rom us can be determined through a combination of (a) studying the light curve. (b) determining the source size. (c) measuring (he astrometric effects of lensing. and (d) detecting the lens or placing limits on the flux we receive [rom it.," We show that in many cases the lens mass and its distance from us can be determined through a combination of (a) studying the light curve, (b) determining the source size, (c) measuring the astrometric effects of lensing, and (d) detecting the lens or placing limits on the flux we receive from it."
+ These same procedures can determine if à planet orbits a star ancl. if it does. can measure (lie mass ratio and projected orbital separation.," These same procedures can determine if a planet orbits a star and, if it does, can measure the mass ratio and projected orbital separation."
+ Such a rich set of tests is available that independent measurenientis of key quantities. such as the lens mass. can be made in some cases.," Such a rich set of tests is available that independent measurements of key quantities, such as the lens mass, can be made in some cases."
+ In 4. we explicitly consider the case in which the lens is a free-floating planet and demonstrate that. especially for nearby planets. the model can be well tested and that mass measurements may be possible in some cases.," In 4, we explicitly consider the case in which the lens is a free-floating planet and demonstrate that, especially for nearby planets, the model can be well tested and that mass measurements may be possible in some cases."
+ 85 sketches (he advantages of focusing attention on events of short duration., 5 sketches the advantages of focusing attention on events of short duration.
+" If the mass of the lens is AJ ancl its distance Irom us is D,. (hen the Einstein anele is: lere. c=D,/Ds. and Dy is the distance to (the lensed source."," If the mass of the lens is $M$ and its distance from us is $D_L,$ then the Einstein angle is: Here, $x=D_L/D_S,$ and $D_S$ is the distance to the lensed source."
+ When the angular separation between the source and lens is 65; (205. 3.505). the magnification is 34% (6%. 1*4).," When the angular separation between the source and lens is $\theta_E$ $2\, \theta_E$, $3.5\, \theta_E$ ), the magnification is $34\%$ $6\%$, $1\%$ )."
+ The relative angular speed between source and lens is w. For nearby lenses. we have For lenses closer to the lensed source. the relative motion may not be dominated by the motion of the lens.," The relative angular speed between source and lens is $\omega.$ For nearby lenses, we have For lenses closer to the lensed source, the relative motion may not be dominated by the motion of the lens."
+ The expression for w is more properly represented. as a sum of terms. but we can use Equation 2 to derive an approximate value by setting e equal to the relative (ransverse speed.," The expression for $\omega$ is more properly represented as a sum of terms, but we can use Equation 2 to derive an approximate value by setting $v$ equal to the relative transverse speed."
+The optimal time delay is the one that minimizes the reduced Uu between the model and the data points.,The optimal time delay is the one that minimizes the reduced $\chi_{red}^{2}$ between the model and the data points.
+ It is important to insist on the difference between y as defined in Eq., It is important to insist on the difference between $\chi^{2}$ as defined in Eq.
+ 2 and on the other hand the reduced Uu in which y is divided by the number of data points in between the light curves of the quasar images for a given time delay., \ref{eqn:chi2} and on the other hand the reduced $\chi_{red}^{2}$ in which $\chi^{2}$ is divided by the number of data points in between the light curves of the quasar images for a given time delay.
+ Indeed. the longer the time delay one tests. the fewer points these light curves have in common. which tends to reduce the y and result in a bias towards longer time delays. hence the use of Uu to avoid this bias.," Indeed, the longer the time delay one tests, the fewer points these light curves have in common, which tends to reduce the $\chi^{2}$ and result in a bias towards longer time delays, hence the use of $\chi_{red}^{2}$ to avoid this bias."
+ A second important difference from the original version of the method is technical: for computational reasons. the length of the model curve should be a power of two. which in some cases proves to be too long it comparison to the data. thus falsifying the balance between data and smoothing terms.," A second important difference from the original version of the method is technical: for computational reasons, the length of the model curve should be a power of two, which in some cases proves to be too long in comparison to the data, thus falsifying the balance between data and smoothing terms."
+ In the original version the smoothing term was applied to the full length of the model., In the original version the smoothing term was applied to the full length of the model.
+ We adapted the program in such a way that the part of the model that is only needed to complete the length until the next power of two. is no longer taken into account in the minimization process.," We adapted the program in such a way that the part of the model that is only needed to complete the length until the next power of two, is no longer taken into account in the minimization process."
+ In this way. the nethod becomes independent of the number of data points i1 the light curve. which was not the case in its original form.," In this way, the method becomes independent of the number of data points in the light curve, which was not the case in its original form."
+ This method has not only been implemented for two light curves. but also for deriving time delays from three and four light curvessimultaneouslv.," This method has not only been implemented for two light curves, but also for deriving time delays from three and four light curves."
+. The strength of this simultaneous approach lies not only in the improved constraints on. the model. but also in that we assume the coherence betweer pairs of time delays. differences in magnitudes. and the slope parameter values.," The strength of this simultaneous approach lies not only in the improved constraints on the model, but also in that we assume the coherence between pairs of time delays, differences in magnitudes, and the slope parameter values."
+ The robustness of the measured time delay is tested i two ways., The robustness of the measured time delay is tested in two ways.
+ First of all. we iteratively attempt to find the three parameters of the model light curve: the spacing of the model curve's sampling points. the range of the smoothing term. and the Lagrange multiplier.," First of all, we iteratively attempt to find the three parameters of the model light curve: the spacing of the model curve's sampling points, the range of the smoothing term, and the Lagrange multiplier."
+ The results should be independent of these parameters as long as we remain in a certain range adapted to the data., The results should be independent of these parameters as long as we remain in a certain range adapted to the data.
+ In a second step. we wish to test the influence of each individual point of the light curve on the time delay.," In a second step, we wish to test the influence of each individual point of the light curve on the time delay."
+ This i5 achieved by means of a classical Jackknife test: for a light curve consisting of data points. we recalculate times the time delay in the light curve of N-I data by successively leaving out one data point at a time.," This is achieved by means of a classical jackknife test: for a light curve consisting of N data points, we recalculate N times the time delay in the light curve of N-1 data by successively leaving out one data point at a time."
+ Time delays should not change drastically because of the removal of à single point from the light curve., Time delays should not change drastically because of the removal of a single point from the light curve.
+ If they do. we know which data point is responsible for the change and we can have a closer look at it.," If they do, we know which data point is responsible for the change and we can have a closer look at it."
+ Errors are calculated by means of Monte Carlo simulations., Errors are calculated by means of Monte Carlo simulations.
+ Normally distributed random errors with the appropriate standard deviation are added to the model light curve and the time delay is redetermined., Normally distributed random errors with the appropriate standard deviation are added to the model light curve and the time delay is redetermined.
+ We note that errors are not added to the data as they already contain the observed error. so adding another would bias the results.," We note that errors are not added to the data as they already contain the observed error, so adding another would bias the results."
+ The model. to which the measurement errors are added. is assumed to provide a more accurate description of the real light curve of the quasar than the data.," The model, to which the measurement errors are added, is assumed to provide a more accurate description of the real light curve of the quasar than the data."
+ This procedure is repeated at least 1000 times. preferably on different combinations of smoothing parameters.," This procedure is repeated at least 1000 times, preferably on different combinations of smoothing parameters."
+ The mean value of the time delay distribution that we obtain is considered to be the final time delay and its dispersion represents the ] «c error bar., The mean value of the time delay distribution that we obtain is considered to be the final time delay and its dispersion represents the 1 $\sigma$ error bar.
+ When we have a markedly asymmetrical distribution. we take its mode as the final time delay and use the 68% confidence intervals to obtain error bars.," When we have a markedly asymmetrical distribution, we take its mode as the final time delay and use the $68\%$ confidence intervals to obtain error bars."
+ In this paper. all quoted uncertainties are | o error bars except where mentioned explicitly.," In this paper, all quoted uncertainties are 1 $\sigma$ error bars except where mentioned explicitly."
+ The advantages of this method are manifold., The advantages of this method are manifold.
+ First. none of the light curves is taken as a reference curve: they are all treated on an equal basis.," First, none of the light curves is taken as a reference curve; they are all treated on an equal basis."
+ Second. a model light curve is obtained for the intrinsic. variations in the quasar. which is also the case for the polynomial fit method described by ?.. but not for the minimum dispersion method developed by ?..," Second, a model light curve is obtained for the intrinsic variations in the quasar, which is also the case for the polynomial fit method described by \citet{2006ApJ...640...47K}, but not for the minimum dispersion method developed by \citet{1996AA...305...97P}."
+ This is important when calculating the error bars. thus avoiding adding random errors to the data.," This is important when calculating the error bars, thus avoiding adding random errors to the data."
+ Finally. since the model is purely numerical. no assumption is made about the quasar's intrinsic light curve. except that it is sufficiently smooth. and we only interpolate the model. never the data.," Finally, since the model is purely numerical, no assumption is made about the quasar's intrinsic light curve, except that it is sufficiently smooth, and we only interpolate the model, never the data."
+ The second method we use is derived from the minimum dispersion method by ? with a number of adjustments as described in. 2.., The second method we use is derived from the minimum dispersion method by \citet{1996AA...305...97P} with a number of adjustments as described in \citet{2010arXiv1009.1473C}.
+ The main improvements to the original ? method consist in: Since no model light curve is constructed. computation time is a lot shorter than for the NMF method.," The main improvements to the original \citet{1996AA...305...97P} method consist in: Since no model light curve is constructed, computation time is a lot shorter than for the NMF method."
+ By using two methods based on completely different principles. we are able to check whether the derived time delays are independent of the method. thus testing their robustness (1.e. independence of the particular way in which the data are analysed).," By using two methods based on completely different principles, we are able to check whether the derived time delays are independent of the method, thus testing their robustness (i.e. independence of the particular way in which the data are analysed)."
+ We now present the main results of our time delay analysis for each of the published light curves of 11 gravitationally lensed quasars., We now present the main results of our time delay analysis for each of the published light curves of 11 gravitationally lensed quasars.
+The magnitude of the luminosity is similar to that produced by the initial breakout radiation. as is seen in numerical simulations (Grasberg2010).,"The magnitude of the luminosity is similar to that produced by the initial breakout radiation, as is seen in numerical simulations \citep{grasberg87, moriya10}."
+". The luminosity lasts until the shock wave al 2 reaches the edge of the dense wind. Rye at lye=LLBSPOMN?DUORE vy, where Rays is in units of LOM em."," The luminosity lasts until the shock wave at $R$ reaches the edge of the dense wind, $R_w$ , at $t_w=1.1E_{51}^{-0.5}M_{e1}^{0.25}D_*^{0.25}R_{w16}^{1.25}$ yr, where $R_{w16}$ is in units of $10^{16}$ cm."
+ Fig., Fig.
+ La illustrates ihe luminosity evolution with the late flattening from the shell interaction., 1a illustrates the luminosity evolution with the late flattening from the shell interaction.
+ In the case of shock breakout [rom a red supergiant. (he shock front takes ~1 day to ivaverse the star and the time for shock breakout is ~LO’ s (IXlein&Chevalier1973).," In the case of shock breakout from a red supergiant, the shock front takes $\sim 1$ day to traverse the star and the time for shock breakout is $\sim 10^3$ s \citep{klein78}."
+. The shock breakout timescale is much less (han the time since explosion., The shock breakout timescale is much less than the time since explosion.
+ In the dense mass loss case considered here. (he Gime for the shock front to move to the breakout region is eRy/esp. which is also the timescale for the breakout event.," In the dense mass loss case considered here, the time for the shock front to move to the breakout region is $\sim R_d/v_{sh}$, which is also the timescale for the breakout event."
+ This property of the huminosity evolution can be seen in simulations of such events (Crassbergetal.1971:Falk&Arnett1977:Chugaietal.2004:Moriva 2010).," This property of the luminosity evolution can be seen in simulations of such events \citep{grassberg71,falk77,chugai04a,moriya10}."
+. The rise to maximum light can have complications due {o variations in (he gas opacity., The rise to maximum light can have complications due to variations in the gas opacity.
+ Al the high cireumstellar densities considered here. it is likely that the gas is initially neutral and that most of the opacity is due to dust in the presupernova environment.," At the high circumstellar densities considered here, it is likely that the gas is initially neutral and that most of the opacity is due to dust in the presupernova environment."
+ The radiation dominated shock wave from the supernova has a precursor in the mass loss region that is expected to heat the cireumstellar dust., The radiation dominated shock wave from the supernova has a precursor in the mass loss region that is expected to heat the circumstellar dust.
+ As the temperature rises through. 1000—2000 Ix. the dust evaporates. giving a decrease in the opacily and the photospheric radius drops to where there is a sharp gradient in the opacity as the gas becomes ionized.," As the temperature rises through $1000-2000$ K, the dust evaporates, giving a decrease in the opacity and the photospheric radius drops to where there is a sharp gradient in the opacity as the gas becomes ionized."
+ In. Type IP supernovae. this property of the opacity causes a recombination wave (o back into the expanding envelope with a constant. photospheric temperature Z5000—6000 Ix. Ilere. the process is inverted and the photosphere is expected to follow the ionization wave moving out through the dense circumstellar gas.," In Type IIP supernovae, this property of the opacity causes a recombination wave to back into the expanding envelope with a constant photospheric temperature $T\sim 5000-6000$ K. Here, the process is inverted and the photosphere is expected to follow the ionization wave moving out through the dense circumstellar gas."
+ This phase of approximately constant temperature can be seen in simulations (Grassberg 1971)., This phase of approximately constant temperature can be seen in simulations \citep{grassberg71}.
+. Once (he circunmstellar mass is ionized. the photospheric expansion slows and the temperature rises.," Once the circumstellar mass is ionized, the photospheric expansion slows and the temperature rises."
+ The light curve rises fairly sharply due to the temperature rise until the maximum temperature. and luminosity. is reached.," The light curve rises fairly sharply due to the temperature rise until the maximum temperature, and luminosity, is reached."
+" We now consider the case that /2,,« AR.", We now consider the case that $R_w<R_d$ .
+ The shock breakout process begins al Hae., The shock breakout process begins at $t_b\approx R_w/v_{sh}$ .
+" The rise time for (he lisht. curve is J,2δν. where df is the distance in from 4&4, across which the diffusion (ime equals the shock crossing (me."," The rise time for the light curve is $t_r\approx\delta R/v_{sh}$, where $\delta R$ is the distance in from $R_w$ across which the diffusion time equals the shock crossing time."
+" We (hus have (ORPRDR,2/6=olt/va. leading to ly2(fva)Ry)."," We thus have $(\delta R)^2\kappa DR_w^{-2}/c=\delta R/v_{sh}$, leading to $t_{r}\approx (R_w/v_{sh})(R_w/R_d)$."
+ In this case the rise time can be considerably shorter that the Gime for initial heating of (he envelope (Fie., In this case the rise time can be considerably shorter that the time for initial heating of the envelope (Fig.
+ Ib)., 1b).
+" Since OR/R,<1. the region involved in shock breakout is not racially extended. so the pressure of (he escaping radiation can accelerate the gas out to τω (Ensiman1994)."," Since $\delta R/R_w<1$, the region involved in shock breakout is not radially extended, so the pressure of the escaping radiation can accelerate the gas out to $R_w$ \citep{ensman94}."
+. Hieh gas velocities are expected around the time of maximum luminosity., High gas velocities are expected around the time of maximum luminosity.
+ A dense shell forms when the radiation can escape. butdoes not produce continuing high Iuminosityv because (he densegas does nol extend far beyond the breakout point and it has been radiativelv accelerated.," A dense shell forms when the radiation can escape, butdoes not produce continuing high luminosity because the densegas does not extend far beyond the breakout point and it has been radiatively accelerated."
+iis the closest super-massive black hole (SALBLD). naturally. residing in the centre of our Galaxy (6.8...Schódeletal.2002:GhezetaL. 2003).. and weighing about Alpy~3.5 M...,"is the closest super-massive black hole (SMBH), naturally residing in the centre of our Galaxy \citep[e.g., ][]{Schoedel02,Ghez03b}, and weighing about $M_{\rm BH} \sim 3.5 \times 10^6 \msun$ ."
+ Pew tens of close voung massive Πο) emission ine stars dominate the energy output of the central parsec of he Galaxy (Ixrabbeetal.1995).," Few tens of close young massive ""He-I"" emission line stars dominate the energy output of the central parsec of the Galaxy \citep{Krabbe95}."
+.. Almost all of these stars xiong to one of twoο rings (Levin&Beloborocoy.2003:CGenzelctal.. 2003)..," Almost all of these stars belong to one of two rings \citep{Levin03,Genzel03a}."
+" The ages of the voung stars are estimated at /c441 Million. vears (CGenzel 2005. o»ivate communication). and the initial masses of the ] stars are probably larger than A240M. (Ixrabbectal..1995:Paumardetabl. 2001)... """," The ages of the young stars are estimated at $t \approx 4\pm 1$ Million years (Genzel 2005, private communication), and the initial masses of the He-I stars are probably larger than $M\simgt 40 \msun$ \citep{Krabbe95,Paumard01}. “"
+"Standard modes of star ormation at 2=0.1/4 pc distances from a SMDIL are orbidden due to a huge tidal field of the central object as the required gas density is ngx1l0ttem2,1.","Standard” modes of star formation at $R= 0.1R_{0.1}$ pc distances from a SMBH are forbidden due to a huge tidal field of the central object as the required gas density is $n_H 
+\simgt 10^{11} \hbox{cm}^{-3} R_{0.1}^{-3}$."
+ Understanding formation mechanisms of these stars is a challenge and a test for both star formation theory anc AGN phenomenon as such., Understanding formation mechanisms of these stars is a challenge and a test for both star formation theory and AGN phenomenon as such.
+ Two models suggested to resolve this dillicultv clo so in two opposite wavs., Two models suggested to resolve this difficulty do so in two opposite ways.
+ CGorhard.(2001). showed that the stars may have been formed at distance of tens of parsecs. ius avoiding the need for excessive gas density prior to star formation. in a massive star cluster.," \cite{Gerhard01} showed that the stars may have been formed at distance of tens of parsecs, thus avoiding the need for excessive gas density prior to star formation, in a massive star cluster."
+ Phe star clusters orbit would then decay through dynamical friction with the background stars., The star cluster's orbit would then decay through dynamical friction with the background stars.
+" Further detailed models however showed jit the cluster needs to be very massive CM~10"" M.) aad very compact (Ixim&Morris.2003:HansenMilosayljevic.Gürkan&Rasio. 2005)."," Further detailed models however showed that the cluster needs to be very massive $M \sim 10^6\msun$ ) and very compact \citep{Kim03,Hansen03,McMillan03,Kim04,Gurkan05}."
+. The other model (Levin&Beloborodov.2003:MilosavlI-jevié&Loeb.2004:NavakshinC'uadra.2005). overcomes the tidal density limit bv sheer force (mass). which naturally occurs at large distances in a standard aceretion disk mocel (Shakura&Sunyaev.1973). when the disk mass exceeds a," The other model \citep{Levin03,Milosavljevic04,NC05} overcomes the tidal density limit by sheer force (mass), which naturally occurs at large distances in a standard accretion disk model \citep{Shakura73} when the disk mass exceeds a"
+expect such objects to have a range of different metallicities with galaxies with significant observed. dust. obseuration such as HICMG6À (2=6.56.2)) likely to be more chemically enriched.,expect such objects to have a range of different metallicities with galaxies with significant observed dust obscuration such as HCM6A $z=6.56$ \citet{Chary:05}) ) likely to be more chemically enriched.
+ The presence of many massive stars in (hese svstems will also result in hieh FUV radiation field strengths and cosmic rav ionisation rates., The presence of many massive stars in these systems will also result in high FUV radiation field strengths and cosmic ray ionisation rates.
+ Collapse is likely to occur al moderately hieh densiües due to the lack of available molecular coolants compared to a Milkv. Way type object., Collapse is likely to occur at moderately high densities due to the lack of available molecular coolants compared to a Milky Way type object.
+" In order to model such a svstem. we calculate the star formation (imescales for (hree different models - (1) n=10 7. €=0.058... €=10PF ! and y=L7x107 . (IL) n10 *.£—0.05... C—I0Hs Land y—L7x10* and (HII) p»=10 ὃν =E.. C—10I ! and 4=L7x10° - at depths corresponding to A,~3 and sl.~10."," In order to model such a system, we calculate the star formation timescales for three different models - (I) $n=10^4$ $^{-3}$, $\xi=0.05\xi_{\odot}$ , $\zeta=10^{-13}$ $^{-1}$ and $\chi=1.7\times10^4$ (II) $n=10^5$ $^{-3}$, $\xi=0.05\xi_{\odot}$, $\zeta=10^{-14}$ $^{-1}$ and $\chi=1.7\times10^3$ and (III) $n=10^5$ $^{-3}$, $\xi=\xi_{\odot}$, $\zeta=10^{-14}$ $^{-1}$ and $\chi=1.7\times10^3$ - at depths corresponding to $A_v\sim3$ and $A_v\sim10$."
+ As previously stated. the presence of many massive stars in hieh redshift svstems justifies our choice of elevated FUV radiation field strengths when modelling such svstems.," As previously stated, the presence of many massive stars in high redshift systems justifies our choice of elevated FUV radiation field strengths when modelling such systems."
+ In order to mimic starburst galaxies with high cosmic rav [hixes such as M82 (2). and NGC253 (2). à hieh redshift. we adopt high cosmic ray ionisation rates.," In order to mimic starburst galaxies with high cosmic ray fluxes such as M82 \citep{Suchkov:m82cr} and NGC253 \citep{Bradford:n253cr} at high redshift, we adopt high cosmic ray ionisation rates."
+ ? estimate a cosmic rav lonization rale of 5xLOPs F for a star formation rate of &100M.. 1 and we expect high redshift ealaxies with higher star lormation rates to have elevated cosmic rav. fluxes., \citet{Meijerink:cr} estimate a cosmic ray ionization rate of $5\times10^{-15}$ $^{-1}$ for a star formation rate of $\sim$ $_\odot$ $^{-1}$ and we expect high redshift galaxies with higher star formation rates to have elevated cosmic ray fluxes.
+ Finally. there is evidence (hat high redshift galaxies exist al a range of metallicities with galaxies with sienificant observed dust obscuration likely to be more chemically enriched and we therefore consider both solar and sub-solar metallicities.," Finally, there is evidence that high redshift galaxies exist at a range of metallicities with galaxies with significant observed dust obscuration likely to be more chemically enriched and we therefore consider both solar and sub-solar metallicities."
+ Our results are sumnmarised in Table 1.., Our results are summarised in Table \ref{tab:highz}. .
+" We can see that for the first model. (he free-fall time is comparable to the cooling time and cores can collapse to form stars even at depths as low as A,~3."," We can see that for the first model, the free-fall time is comparable to the cooling time and cores can collapse to form stars even at depths as low as $A_v\sim3$."
+ The formation of icv mantles is likely to be inhibited in such systenis as the Ireeze-out (nie is very. large compared to the desorption time., The formation of icy mantles is likely to be inhibited in such systems as the freeze-out time is very large compared to the desorption time.
+ The second model has slightly lower cosmic ray f[Inxes and FUY radiation field strengths but is considered (ο be at a higher density., The second model has slightly lower cosmic ray fluxes and FUV radiation field strengths but is considered to be at a higher density.
+" In this case. the free-fall timescale is shorter than the cooling Gime at low 21, and cores can collapse to form stars but at vl.10. the cooling lime is larger ancl (he propensity for low mass star formation is reduced."," In this case, the free-fall timescale is shorter than the cooling time at low $A_v$ and cores can collapse to form stars but at $A_v\sim10$, the cooling time is larger and the propensity for low mass star formation is reduced."
+ Once again. the formation of 1ον mantles is inhibited in such a svstem.," Once again, the formation of icy mantles is inhibited in such a system."
+ The third model also has elevated FUV radiation field strengths and cosmic ray [Inxes but is now considered at solar metallicity., The third model also has elevated FUV radiation field strengths and cosmic ray fluxes but is now considered at solar metallicity.
+ The cooling time is now significantly lower than (he free-fall time at all depths ancl low-mass star formation is unimpeded by thermal pressure while ice formation still remains inhibited., The cooling time is now significantly lower than the free-fall time at all depths and low-mass star formation is unimpeded by thermal pressure while ice formation still remains inhibited.
+ In all three cases. the ambipolar diffusion timescale is very huge so the residual field mav be able to halt collapse.," In all three cases, the ambipolar diffusion timescale is very large so the residual field may be able to halt collapse."
+ Using the temperaturesand densities of (hese model runs.we calculate the thermal ancl magnetic pressure for each of themodels in the same wav as in," Using the temperaturesand densities of these model runs,we calculate the thermal and magnetic pressure for each of themodels in the same way as in"
+us with a possibility to be independent from symmetry assumptions and to obtain more realistic deprojection temperature maps.,us with a possibility to be independent from symmetry assumptions and to obtain more realistic deprojection temperature maps.
+ Shock fronts are places of gas heating and. therefore. the gas temperature inhomogeneities along the line-of-sight are large near shock fronts.," Shock fronts are places of gas heating and, therefore, the gas temperature inhomogeneities along the line-of-sight are large near shock fronts."
+ Studving the temperature variance map provides us with a new independent approach to unveil shock fronts in merging galaxy clusters., Studying the temperature variance map provides us with a new independent approach to unveil shock fronts in merging galaxy clusters.
+ We now demonstrate how to derive the temperature variance along the line-of-sight for hot galaxy clusters with gas temperatures in the range of 10 keV — 15 keV by using the improved method proposed in the previous section., We now demonstrate how to derive the temperature variance along the line-of-sight for hot galaxy clusters with gas temperatures in the range of 10 keV – 15 keV by using the improved method proposed in the previous section.
+ To produce the SZ intensity maps at frequencies of 150 GHz. 197 GHz. and 397 GHz. we use the 3D numerical simulations of the merging hot galaxy cluster presented in Akahori Yoshikawa (2010).," To produce the SZ intensity maps at frequencies of 150 GHz, 197 GHz, and 397 GHz, we use the 3D numerical simulations of the merging hot galaxy cluster presented in Akahori Yoshikawa (2010)."
+ The simulations of merging galaxy clusters by Akahori Yoshikawa (2010) are carried out using an N-body and SPH code developed in Akahort Yoshikawa (2008). which adopts the entropy—conservative formulation of SPH by Springel Hernquist (2002). and the standard Monaghan-Gingold artificial viscosity (Monaghan Gingold 1983) with the Balsara limiter (Balsara 1995).," The simulations of merging galaxy clusters by Akahori Yoshikawa (2010) are carried out using an N-body and SPH code developed in Akahori Yoshikawa (2008), which adopts the entropy–conservative formulation of SPH by Springel Hernquist (2002), and the standard Monaghan-Gingold artificial viscosity (Monaghan Gingold 1983) with the Balsara limiter (Balsara 1995)."
+ They carried out an adiabatic run in which radiative cooling. star formation. and AGN feed back are not taken into consideration.," They carried out an adiabatic run in which radiative cooling, star formation, and AGN feed back are not taken into consideration."
+ In the simulations. an encounter of two free-falling galaxy subclusters from the epoch of their turn-around is considered.," In the simulations, an encounter of two free-falling galaxy subclusters from the epoch of their turn-around is considered."
+" The two galaxy subelusters have virial masses and radii of M94,= and roo.=191 Mpe. and Afi»=2x107 and rope=M.01.231 Mpe. respectively. where roy) is the radius M.within which the mean cluster mass density is 200 times the present cosmic critical density."," The two galaxy subclusters have virial masses and radii of $M_{\mathrm{200},1}=8\times10^{14} M_\odot$ and $r_{\mathrm{200},1}=1.91$ Mpc, and $M_{\mathrm{200},2}=2\times10^{14} M_\odot$ and $r_{\mathrm{200},2}=1.21$ Mpc, respectively, where $r_{\rm 200}$ is the radius within which the mean cluster mass density is 200 times the present cosmic critical density."
+ The initial distributions of a dark matter halo and à ICM component are the NFW density profile (Navarro et al., The initial distributions of a dark matter halo and a ICM component are the NFW density profile (Navarro et al.
+ 1997) and the £-model density profile (Cavaliere Fusco-Femiano 1976). respectively.," 1997) and the $\beta$ -model density profile (Cavaliere Fusco-Femiano 1976), respectively."
+" The scale radius of the NFW profile is set to r,=r:99/5.16. and the core radius of the 6-model is r.=r,/2."," The scale radius of the NFW profile is set to $r_{\rm s} = r_{\rm 200}/5.16$, and the core radius of the $\beta$ -model is $r_{\rm c}=r_{\rm s}/2$."
+ The massive subcluster is composed of one million dark matter particles and the same number of SPH particles. while he less massive subcluster is represented by a quarter million yarticles for each component.," The massive subcluster is composed of one million dark matter particles and the same number of SPH particles, while the less massive subcluster is represented by a quarter million particles for each component."
+ The corresponding spatial resolution (the smoothing length) of SPH is ~560 kpe at the ICM density of 10em.," The corresponding spatial resolution (the smoothing length) of SPH is $\sim 50$ kpc at the ICM density of $10^{-3} {\rm
+cm^{-3}}$."
+ We use the result with zero impact parameter at a ime of/=0.5 Gyr. where ¢=0 Gyr corresponds to the time of he closest approach of the centers of the dark matter halos.," We use the result with zero impact parameter at a time of $t=0.5$ Gyr, where $t=0$ Gyr corresponds to the time of the closest approach of the centers of the dark matter halos."
+ The average temperature of this simulated galaxy cluster is = 10 keV and the temperature of the hottest regions associated with shock τοηῖς is = [5 keV. For further details of the simulation. see Akahori Yoshikawa (2010).," The average temperature of this simulated galaxy cluster is $\simeq$ 10 keV and the temperature of the hottest regions associated with shock fronts is $\simeq$ 15 keV. For further details of the simulation, see Akahori Yoshikawa (2010)."
+ The dense cores of these two subclusters are accelerated by he preceding dark matter halos and shock waves are formed in ront of the cores., The dense cores of these two subclusters are accelerated by the preceding dark matter halos and shock waves are formed in front of the cores.
+ Shock fronts are shown in the Mach number distribution plotted in Fig., Shock fronts are shown in the Mach number distribution plotted in Fig.
+ 2 of Akahori Yoshikawa (2010)., 2 of Akahori Yoshikawa (2010).
+ Since he temperature and density maps of the simulated hot merging cluster are inhomogeneous owing to the presence of shock waves. we study the map of the temperature variance along the sight for this cluster.," Since the temperature and density maps of the simulated hot merging cluster are inhomogeneous owing to the presence of shock waves, we study the map of the temperature variance along the line-of-sight for this cluster."
+" For a hot galaxy cluster with inhomogeneous density and temperature (&,7..€1S keV) distributions. the CMB intensity change produced by the SZ effect in the formalism based on an extension of the Kompaneets equation. which is correct to second order in the expansion parameter ©. can be written as where οἱ). ο). and ον) are the spectral functions taken from Challinor Lasenby (1998)."," For a hot galaxy cluster with inhomogeneous density and temperature $k_{\rm{b}} T_{\rm{e}}\leq 15$ keV) distributions, the CMB intensity change produced by the SZ effect in the formalism based on an extension of the Kompaneets equation, which is correct to second order in the expansion parameter $\Theta$, can be written as where $g(x)$, $g_{1}(x)$, and $g_{2}(x)$ are the spectral functions taken from Challinor Lasenby (1998)."
+ Note that higher order terms should be included for calculating the SZ ettect from very hot galaxy clusters with KZ.21S keV (e.g. Challinor Lasenby 1998: Itoh et al., Note that higher order terms should be included for calculating the SZ effect from very hot galaxy clusters with $k_{\rm{b}} T_{\rm{e}}\geq 15$ keV (e.g. Challinor Lasenby 1998; Itoh et al.
+ 1998)., 1998).
+ Since the optical depth 7 of the electron plasma can be derived from X-ray or SZ observations. we rewrite Eq. (12))," Since the optical depth $\tau$ of the electron plasma can be derived from X-ray or SZ observations, we rewrite Eq. \ref{form22}) )"
+" as In terms of the SZ intensities. the temperature variance along the line-of-sight is given by where Aii is the SZ intensity at a frequency of 150 GHz (x=2.64). ΔΙτοι, is the combined SZ intensity at frequencies of 197 GHz (x=3.47) and 397 GHz (x=7.0). which is detined as (Prokhorov et al."," as In terms of the SZ intensities, the temperature variance along the line-of-sight is given by where $\Delta I_{\mathrm{150GHz}}$ is the SZ intensity at a frequency of 150 GHz (x=2.64), $\Delta \tilde{I}_{\mathrm{197GHz},
+\mathrm{397GHz}}$ is the combined SZ intensity at frequencies of 197 GHz (x=3.47) and 397 GHz (x=7.0), which is defined as (Prokhorov et al."
+ 20100) and 2 is the coethcient which equals The map of the standard temperature deviation along the line-of-sight derived from our mock SZ observations at frequencies of 150 GHz. 197 GHz. and 397 GHz is shown in Fig.," 2010a) and $\tilde{g}$ is the coefficient which equals The map of the standard temperature deviation along the line-of-sight derived from our mock SZ observations at frequencies of 150 GHz, 197 GHz, and 397 GHz is shown in Fig."
+ 3 (top panel)., \ref{Fig3} (top panel).
+ The map of the temperature variance. calculated by using the simulated 3D gas temperature and density distributions instead of the SZ intensities at frequencies of 150 GHz. 197 GHz. and 397 GHz. is plotted in Fig.," The map of the temperature variance, calculated by using the simulated 3D gas temperature and density distributions instead of the SZ intensities at frequencies of 150 GHz, 197 GHz, and 397 GHz, is plotted in Fig."
+ 3. (bottom panel)., \ref{Fig3} (bottom panel).
+ Comparing the maps in Fig., Comparing the maps in Fig.
+ 3. shows that the maximal ditference between these maps is =25%.., \ref{Fig3} shows that the maximal difference between these maps is $\approx$.
+ Measurements of the temperature variance by means of the proposed method permit us to study temperature inhomogeneities along the line-of-sight., Measurements of the temperature variance by means of the proposed method permit us to study temperature inhomogeneities along the line-of-sight.
+ We find that the values of the the temperature variance are high in the post-shock regions and. therefore. the map of this quantity can also be used to unveil merger shocks in hot galaxy clusters.," We find that the values of the the temperature variance are high in the post-shock regions and, therefore, the map of this quantity can also be used to unveil merger shocks in hot galaxy clusters."
+ We will study possibilities of deriving the temperature variance along the line-of-sight for hot merging galaxy clusters by means of the planned SZ experiments in a forthcoming paper., We will study possibilities of deriving the temperature variance along the line-of-sight for hot merging galaxy clusters by means of the planned SZ experiments in a forthcoming paper.
+ A number of important gasdynamic processes can be studied by observing merging galaxy clusters., A number of important gasdynamic processes can be studied by observing merging galaxy clusters.
+ Studying the temperature structure of merging clusters is important to reveal heated regions associated with merger shock waves and is complicated because the temperature maps of merging clusters are very inhomogeneous and because the approximation of a spherically symmetric gas distribution is not suitable for merging galaxy clusters., Studying the temperature structure of merging clusters is important to reveal heated regions associated with merger shock waves and is complicated because the temperature maps of merging clusters are very inhomogeneous and because the approximation of a spherically symmetric gas distribution is not suitable for merging galaxy clusters.
+ In many merging clusters. such as the Abell 3376 cluster (see Bagchi et al.," In many merging clusters, such as the Abell 3376 cluster (see Bagchi et al."
+It is well accepted that centimeter wavebauc clnission frou ACINS is associated with a jet of svuchotron plasma. the accretion structure ac iuniediate environment of a supermassive black hole coutributing broad band enmüssiou frou the infrared to the eauuna ray spectrun.,"It is well accepted that centimeter waveband emission from AGNs is associated with a jet of synchotron plasma, the accretion structure and immediate environment of a supermassive black hole contributing broad band emission from the infrared to the gamma ray spectrum."
+ Tempora various are observed in the radio flux. aux a number of processes have been proposed to explain this. such as an accretion rate that nay change with time. an accretion disk that exhibits instability. an outflow that iuav be IWelvin-UWehuboltz unstable. or an outflow that nay iuteract with ambient inhomoecueitics.," Temporal variations are observed in the radio flux, and a number of processes have been proposed to explain this, such as an accretion rate that may change with time, an accretion disk that exhibits instability, an outflow that may be Kelvin-Helmholtz unstable, or an outflow that may interact with ambient inhomogeneities."
+ The cluporal variations in the radio flux motivate a search for characteristic time scales. aud. if fore. would lend insielit ito the mechanisms causine he variations.," The temporal variations in the radio flux motivate a search for characteristic time scales, and, if found, would lend insight into the mechanisms causing the variations."
+ Iu practice. searching for possible quasiperiodic ychavior masked by a stochastic component can ve rather difficult. especially iu the context of irregular time sampling.," In practice, searching for possible quasiperiodic behavior masked by a stochastic component can be rather difficult, especially in the context of irregular time sampling."
+ One promising technique is to perform a continuous wavelet trausform on the signal. and map out the coefficieuts in wavelet space.," One promising technique is to perform a continuous wavelet transform on the signal, and map out the coefficients in wavelet space."
+ This was done on data for the DL Lac OJ 287. vesulting in evidence for periods iu the radio spectral region of ~ 1.66 vr. and ~ 1.12 vr dominating in the 1980s (IDughes.Aller 1998)..," This was done on data for the BL Lac OJ 287, resulting in evidence for periods in the radio spectral region of $\sim$ 1.66 yr, and $\sim$ 1.12 yr dominating in the 1980s \citep{hug98}. ."
+" The results led Uneles.Aller.&Aller(1998). to propose that the results cau best be explained by a ""shock-in-jet model. iu which the longer periodicity is linked with au otherwise quicsceut jet. aud the shorter with a passing shock-wave,"," The results led \citet{hug98} to propose that the results can best be explained by a “shock-in-jet” model, in which the longer periodicity is linked with an otherwise quiescent jet, and the shorter with a passing shock-wave."
+ The continuous wavelet transform detects quasiperiodic behavior through visual examination of the map over translation and dilation. aud assigus a characteristic time scale dy looking for peaks iu the time-averaged wavelet power (see LA).," The continuous wavelet transform detects quasiperiodic behavior through visual examination of the map over translation and dilation, and assigns a characteristic time scale by looking for peaks in the time-averaged wavelet power (see 4.1)."
+ The scales that the peaks occur at correspond to Fourier periods. from) which a characteristic time scale aay be deduced.," The scales that the peaks occur at correspond to Fourier periods, from which a characteristic time scale may be deduced."
+ Although the continuous transform is effective. when assiguing a characteristic time scale it docs not fully iuclude information of how a timeseries," Although the continuous transform is effective, when assigning a characteristic time scale it does not fully include information of how a timeseries"
+" (Fischeretal.2001:Vogt2000). «7.5. L + « 15701 « 13.5401. à<<207 V7.5. Weuryctal.(1996):Tinneyet loeR4, jo Butleretal.(1996) Butleretal. C2001).. Tn addition to our primarysampleofchwarts. V7."," \citep{fischer01,vmba00}, \citep{AAPSI,AAPSII,AAPSIII,AAPSIV,AAPSV,AAPSVI} $<$ $_2$ $^{-1}$ $\times$ $\mu$ $\times$ $\mu$ $\delta < -20$ $<$ \citet{hsdb96,CaHKI} $\log$ $>$ \citet{bmwmd96} \citet{AAPSII}."
+5V«9) Lanegllin(2000)weby ," $<$$<$ \citet{laughlin00} $uvby$ \citep{AAPSIV,AAPSII}."
+5.08: Tiunevct (Ilouck&Cowley1975).., $-5.08$ \citet{CaHKI}) \citep{mssI}.
+ LS.5E0.8 My L.7240.09 (ESA1997) 113623:0.10 (Cox2000) «eby weby (Schuster&Nissen1989) [Fe/MJ=|0.20+0.16. (Doud2002).," $\pm$ $_V$ $\pm$ \citep{esa97}
+ $\pm$ \citep{allenIV} $uvby$ \citep{hm97}, $uvby$ \citep{sn89} $+0.20{\pm}0.16$ \citep{Bond02},"
+event that occurred in Chis region of the solar svstem: therefore (his sample Likely preserves best the post-nueration orbital distribution of the asteroid belt.,event that occurred in this region of the solar system; therefore this sample likely preserves best the post-migration orbital distribution of the asteroid belt.
+ The proper eccentricity distribution of these asteroids is shown in Fig. 5.., The proper eccentricity distribution of these asteroids is shown in Fig. \ref{f:MBA-big-dist}.
+ This distribution has usually been described in the literature by simply quoting iis mean value (and sometimes a dispersion) (Murray1999:O'Brienοἱal. 2007)..," This distribution has usually been described in the literature by simply quoting its mean value (and sometimes a dispersion) \citep{Murray:1999SSD,OBrien:2007p95}."
+" Our best fit single Gaussian distribution to this data has a mean. jf, and standard deviation. σ,. given by ji.—0.135+0.00013. and 0.07162:0.00022. and is plotted in Fig. 5.."," Our best fit single Gaussian distribution to this data has a mean, $\mu_e$ and standard deviation, $\sigma_e$, given by $\mu_e=0.135\pm0.00013$ and $\sigma_e=0.0716\pm0.00022$ , and is plotted in Fig. \ref{f:MBA-big-dist}."
+ However. we also note (bv eve) a possible indication ola double-peak feature in the observed population.," However, we also note (by eye) a possible indication of a double-peak feature in the observed population."
+ Our best fit double Gaussian* distribution (modelled as (wo symmetrical Gaussians with the same standard deviation. but different mean values) to the same data has the following parameters: Alore details of how these fits were obtained are described in Appendix A.. where we also discuss goodness-of-fit of the single and double gaussian distributions.," Our best fit double Gaussian distribution (modelled as two symmetrical Gaussians with the same standard deviation, but different mean values) to the same data has the following parameters: More details of how these fits were obtained are described in Appendix \ref{sec:sweeprate-appendix-fitting}, where we also discuss goodness-of-fit of the single and double gaussian distributions."
+ A Ixolmogorov-Sinai test shows that there is only a probability that the observed eccentricities actually are Consistent with the best-lit single. Gaussian distribution. and a probability that (μον are consistent with the best-fit double Gaussian distribution: while the double gaussian is apparently a better fit to the data compared to the single eaussian. a ο probability of only is quite far from a statistically significant level of confidence.," A Kolmogorov-Sinai test shows that there is only a probability that the observed eccentricities actually are consistent with the best-fit single Gaussian distribution, and a probability that they are consistent with the best-fit double Gaussian distribution; while the double gaussian is apparently a better fit to the data compared to the single gaussian, a K-S probability of only is quite far from a statistically significant level of confidence."
+ A clip test. [or multi-modality in the observed eccentricity distribution (Ilartigan&Ilartigan1935) is also inconclusive: (he Likely reason lor (his is (hat our data sample is not very large. ancl the separation between (he (wo putative peaks is too small in relation to the clispersion.," A dip test for multi-modality in the observed eccentricity distribution \citep{Hartigan:1985p3924} is also inconclusive; the likely reason for this is that our data sample is not very large, and the separation between the two putative peaks is too small in relation to the dispersion."
+ Nevertheless. we bravely proceed in (he next section with the implications of a eccentricity distribution. with the caveat that some of the conclusions we reach are based on this statistically marginal result.," Nevertheless, we bravely proceed in the next section with the implications of a double-Gaussian eccentricity distribution, with the caveat that some of the conclusions we reach are based on this statistically marginal result."
+ By relating the gj secular Irequency to the seminiajor axis of Saturn. gj can be related to (he migration rate of Saturn. cq.," By relating the $g_6$ secular frequency to the semimajor axis of Saturn, $\dot{g}_6$ can be related to the migration rate of Saturn, $\dot{a}_{6}$."
+ In (his section only the effects of (he sweeping vs secular resonance will be considered. and effects due to sweeping 5.jovian mean motion resonances will be ignored.," In this section only the effects of the sweeping $\nu_6$ secular resonance will be considered, and effects due to sweeping jovian mean motion resonances will be ignored."
+ Because Jupiter is thought to have mieratec inward a much smaller distance (han saturn migrated outward during planetesimal-driven migration. (he effects due to migrating jovian AIAIRs were likely confined (ο narrow regions near strong resonances," Because Jupiter is thought to have migrated inward a much smaller distance than Saturn migrated outward during planetesimal-driven migration, the effects due to migrating jovian MMRs were likely confined to narrow regions near strong resonances"
+21ος of the oscilations. the effect of inultiue accretion streams aud the iuflueunce of the uaenetic field. arguiug he supporting pressure of the magnetic field has to be asou Into accnut.,"period of the oscillations, the effect of multiple accretion streams and the influence of the magnetic field, arguing the supporting pressure of the magnetic field has to be taken into account."
+ Different sudies of time-depevent accretion sliocks COiceuntrate on different aspects., Different studies of time-dependent accretion shocks concentrate on different aspects.
+ ? preseut |D magnueto-ivcdrocvac models of super-Alfvénnic acerction flows with 1ifall veociies around LOO kia 1 and different incident angles., \citet{2008MNRAS.388..357K} present 1D magneto-hydrodynamic models of super-Alfvénnic accretion flows with infall velocities around 100 km $^{-1}$ and different incident angles.
+ In their sinulatious stroig laenetic Ποcadds to a stable accretion flow., In their simulations strong magnetic fields lead to a stable accretion flow.
+ They develop a stability criterion which depends ou the anele between he fied and the stellar surface., They develop a stability criterion which depends on the angle between the field and the stellar surface.
+" However. lis criterion is not been verified for higher iifall veocitics,"," However, this criterion has not been verified for higher infall velocities."
+ If it is still valid for velocities of 525 kll Ἡ1 and densities of. 4LO1?2 & 3I. as measured from. N-rav spectroscopy in TW να (????7).. hen a imagretic field of 100 C; js sufficient to kee) the shocks stabe for an accretion flow perpendicular totιο stellar surace.," If it is still valid for velocities of 525 km $^{-1}$ and densities of $10^{-12}$ g $^{-3}$, as measured from X-ray spectroscopy in TW Hya \citep{2002ApJ...567..434K,twhya,acc_model,2009A&A...505..755R,2010ApJ...710.1835B}, then a magnetic field of 100 G is sufficient to keep the shocks stable for an accretion flow perpendicular to the stellar surface."
+ Yet. it is nof clear if this is also vali for the strneger magnetic fields usually miecasured i1 the accretion sXXt regions of CTTS (HD. which indicate s~Alfvénnic fiows.," Yet, it is not clear if this is also valid for the stronger magnetic fields usually measured in the accretion spot regions of CTTS \citep{2004Ap&SS.292..619V, 2008MNRAS.386.1234D}, which indicate sub-Alfvénnic flows."
+ ο performed 2D simulations of accretion perpendicular to the stellar surface includiug heat conduction i both the supeor-A]vénunic aud the sub-Avéunic regine., \citet{2010A&A...510A..71O} performed 2D simulations of accretion perpendicular to the stellar surface including heat conduction in both the super-Alfv\'ennic and the sub-Alfvénnic regime.
+ Thev coiir the nuportauce of the magetic field. axl fud that shocks wih hie1 J (the ratio heWeel gas pressure aud naenetic pressure) become stable through magnetic danyping for EmrereAfvéónuc flows.," They confirm the importance of the magnetic field, and find that shocks with high $\beta$ (the ratio between gas pressure and magnetic pressure) become stable through magnetic damping for super-Alfvénnic flows."
+ Ux observe quasi-periocic shock oscillations for cases wih luweecr magnetic fieds when the flow becomes sub-Alfvónuice aux chaotic behaviour without periodicity for even weaker maenctic fields., \citet{2010A&A...510A..71O} observe quasi-periodic shock oscillations for cases with larger magnetic fields when the flow becomes sub-Alfvénnic and chaotic behaviour without periodicity for even weaker magnetic fields.
+ ? also compare their results to LD simultions with a similar setup by ? and fined that 2D iodes predict a smaller αλληπαο and higher frequency than 1D models., \citet{2010A&A...510A..71O} also compare their results to 1D simulations with a similar setup by \citet{2008A&A...491L..17S} and find that 2D models predict a smaller amplitude and higher frequency than 1D models.
+ For large values of Jj the magnetic fick cannot confine the accretion stream to a funnel aud iuatter flows out side-waves. ci15urbing the surrounding stellar atiuiospire.," For large values of $\beta$ the magnetic field cannot confine the accretion stream to a funnel and matter flows out side-ways, disturbing the surrounding stellar atmosphere."
+" Altrough 7? and ?— predict stabilitv iu cüffereut reenues, they both agree that there are maeuetie field configuraions that inhibit shock oscilatious iut16 sUupoer-Alfvénnic regine."," Although \citet{2008MNRAS.388..357K} and \citet{2010A&A...510A..71O} predict stability in different regimes, they both agree that there are magnetic field configurations that inhibit shock oscillations in the super-Alfv\'ennic regime."
+ Onlv ? probe sub-zAlfvéóunic flows. but they test less parameters than 7.. so we cannot exclude that siiar stability regions exit for sub-Alfvónuuc flows.," Only \citet{2010A&A...510A..71O} probe sub-Alfvénnic flows, but they test less parameters than \citet{2008MNRAS.388..357K}, so we cannot exclude that similar stability regions exit for sub-Alfvénnic flows."
+ So. we lay not sce periodic oxcillaticuns of the accretion shock. not because of confusion of a larger nunber of accretion streams. but because the shock is intrinsically stable.," So, we may not see periodic oscillations of the accretion shock, not because of confusion of a larger number of accretion streams, but because the shock is intrinsically stable."
+ We preseuted observations of TW να with SALT iu two filters centred just above aud below the Balmer junip with a timing resolution of 0.14 s and observations of AA Tau with the single-photon counting device OPTIMA., We presented observations of TW Hya with SALT in two filters centred just above and below the Balmer jump with a timing resolution of 0.14 s and observations of AA Tau with the single-photon counting device OPTIMA.
+ TW να shows an autocorrelation ou timescales of a few seconds. however we detect no periodic signal in TW IIva or AA Tau.," TW Hya shows an autocorrelation on timescales of a few seconds, however we detect no periodic signal in TW Hya or AA Tau."
+ We place a confidence upper limit on the mused fraction of the lielycurves as 0.001 for TW να in he frequeicy range 0.02-3 IIz in he 310 nni filter aud 1-3 Iz 1i the 280 nm filter., We place a confidence upper limit on the pulsed fraction of the lightcurves as 0.001 for TW Hya in the frequency range 0.02-3 Hz in the 340 nm filter and 0.1-3 Hz in the 380 nm filter.
+ TIi6 correspondius value Or AA Tau is an iuplitide of 0.)05 for 0.02-50 IIz., The corresponding value for AA Tau is an amplitude of 0.005 for 0.02-50 Hz.
+" We sinulate lighteurves to calculate upver Iuuits for scenarios with 1e inclependen accretkn shocks with ciffereut 11ος»,", We simulate lightcurves to calculate upper limits for scenarios with multiple independent accretion shocks with different periods.
+ The timescale seerinthea utocorrelation function is cousistent with the flow time hrough the accretion shock aud he dimension o: affected area as estimated from he amplitudes in the liehteurve., The timescale seen in the autocorrelation function is consistent with the flow time through the accretion shock and the dimension of affected area as estimated from the amplitudes in the lightcurve.
+ Our Imits ou periodicity are so stringent. that we discuss the iiagnoetic field needed ο stabilise the accretio1 shock.," Our limits on periodicity are so stringent, that we discuss the magnetic field needed to stabilise the accretion shock."
+of small Az.,of small $\Delta z$.
+" Hence. we correct the recovered o,(z) using a low-order polynomial fit to the shape of the overall samples JNfdz. but use the fluctuations (compared to a smooth fit) in the observed redshift distribution of the spectroscopic sample dN,/dz. which will be known from the same observations used to perform cross-ecorrelation measurements. to correct for sample variance."," Hence, we correct the recovered $\phi_p(z)$ using a low-order polynomial fit to the shape of the overall sample's $dN/dz$, but use the fluctuations (compared to a smooth fit) in the observed redshift distribution of the spectroscopic sample $dN_s/dz$, which will be known from the same observations used to perform correlation measurements, to correct for sample variance."
+ This correction assumes that deviations from the mean in both samples behave similarly with redshift: we might expect their amplitude to scale with the large-scale-structure bias of a given sample. but we do not apply any correction for that here.," This correction assumes that deviations from the mean in both samples behave similarly with redshift; we might expect their amplitude to scale with the large-scale-structure bias of a given sample, but we do not apply any correction for that here."
+" In tests. we have found that a correction using fluctuations in dN,ας was as effective in constraining parameters as one based on fluctuations in the dNfdz of the overall sample our photometric subsample was selected from. and so we focus on the former. more realistic technique."," In tests, we have found that a correction using fluctuations in $dN_s/dz$ was as effective in constraining parameters as one based on fluctuations in the $dN/dz$ of the overall sample our photometric subsample was selected from, and so we focus on the former, more realistic technique."
+ In more detail. we first divided the recovered distribution by à smooth fit (using a 5th-degree polynomial function) to the overall dN/dz of the entire simulation averaged over all 24 fields.," In more detail, we first divided the recovered distribution by a smooth fit (using a 5th-degree polynomial function) to the overall $dN/dz$ of the entire simulation averaged over all 24 fields."
+ This eliminates gradients associated with the shape of the parent sample's overall redshift distribution. without removing deviations due to sample variance., This eliminates gradients associated with the shape of the parent sample's overall redshift distribution without removing deviations due to sample variance.
+ To correct for the latter. we need to quantify the fluctuations in the spectroscopic sample relative to a mean distribution.," To correct for the latter, we need to quantify the fluctuations in the spectroscopic sample relative to a mean distribution."
+" For this smooth. mean distribution. (4N,/dzj. we used the same fit to the redshift distribution of the spectroscopic sample averaged over all 24 fields which was employed to construct the random catalogs for autocorrelation measurements (83.1)."," For this smooth, mean distribution, $\langle dN_s/dz\rangle$, we used the same fit to the redshift distribution of the spectroscopic sample averaged over all 24 fields which was employed to construct the random catalogs for autocorrelation measurements $\S\ref{sec:autospec}$ )."
+ Using a fit to a given set of four fields would make little difference. as the deviation from the smooth fit at a given redshift bin due to sample variance are much larger than the deviations between the smooth fit to 4 or 24 fields.," Using a fit to a given set of four fields would make little difference, as the deviation from the smooth fit at a given redshift bin due to sample variance are much larger than the deviations between the smooth fit to 4 or 24 fields."
+" We then calculate the ratio dN,dz£/dz). where dN,/dz is the redshift distribution of the spectroscopic sample averaged over the four fields used in that measurement. and correct for sample variance by dividing each measurement of o,(z) by this quantity."," We then calculate the ratio $dN_s/dz/\langle dN_s/dz\rangle$, where $dN_s/dz$ is the redshift distribution of the spectroscopic sample averaged over the four fields used in that measurement, and correct for sample variance by dividing each measurement of $\phi_p(z)$ by this quantity."
+ After applying these corrections to each distribution. each measurement is normalized so that their integral is unity. and then fit for (2) and o. using a normalized Gaussian fitting function.," After applying these corrections to each distribution, each measurement is normalized so that their integral is unity, and then fit for $\langle z \rangle$ and $\sigma_z$ using a normalized Gaussian fitting function."
+ Fig., Fig.
+ 10. shows the median and standard deviation of 10 measurements of the recovered op(z) before and after correcting for sample variance., \ref{fig:phicvar} shows the median and standard deviation of $10^4$ measurements of the recovered $\phi_p(z)$ before and after correcting for sample variance.
+ In both plots the fit to the overall dN/dz is divided out., In both plots the fit to the overall $dN/dz$ is divided out.
+ It is clear to the eye that the distribution corrected for sample variance is a better fit to the underlying selection function; more quantitatively. 1t reduces errors in determining the parameters of the Gaussian selection function by ~10%.," It is clear to the eye that the distribution corrected for sample variance is a better fit to the underlying selection function; more quantitatively, it reduces errors in determining the parameters of the Gaussian selection function by $\sim10\%$."
+ We assess the reconstruction of the photometric sample in two ways., We assess the reconstruction of the photometric sample in two ways.
+ First. we compare the reconstructed parameters. fz) and o-. of the Gaussian selection function to. the true values. known by construction.," First, we compare the reconstructed parameters, $\langle z \rangle$ and $\sigma_z$, of the Gaussian selection function to the true values, known by construction."
+ Second. we compare the reconstructed parameters of the selection function to the parameters of a Gaussian fit to the actual normalized distribution of each set of four fields used.," Second, we compare the reconstructed parameters of the selection function to the parameters of a Gaussian fit to the actual normalized distribution of each set of four fields used."
+ The latter method should be more robust to systematic errors in the πιο dN/dz we divide each measurement by., The latter method should be more robust to systematic errors in the 'true' $dN/dz$ we divide each measurement by.
+" For the first test. where (zjj4,=0.75 and o.,,4,=0.20. we find (Zu—(Sane)=7.796«10707.415 and σσ)=8440«1077E8.545(04. where as usual the values given are the median and standard deviation of all measurements. respectively."," For the first test, where $\langle z\rangle_{true} = 0.75$ and $\sigma_{z,true} = 0.20$ , we find $\langle \langle z \rangle_{rec} - \langle z \rangle_{true}\rangle = 7.796\times10^{-4}\pm7.415\times10^{-3}$ and $\langle \sigma_{z,rec} - \sigma_{z,true} \rangle = 8.140\times10^{-4}\pm8.545\times10^{-3}$, where as usual the values given are the median and standard deviation of all measurements, respectively."
+ The second test. where (Duae aNd σε are determined by a Gaussian fit to the true distribution of each measurement. we find ((24—(Dane)=7.259«10727.465«107 and(624-0244)=4.724.107€8.546« 107.," The second test, where $\langle z\rangle_{true}$ and $\sigma_{z,true}$ are determined by a Gaussian fit to the true distribution of each measurement, we find $\langle \langle z\rangle_{rec}-\langle z\rangle_{true}\rangle = 7.259\times10^{-4}\pm7.465\times10^{-3}$ and$\langle \sigma_{z,rec}-\sigma_{z,true}\rangle = 4.724\times10^{-4}\pm8.546\times10^{-3}$ ."
+ In all cases. the bias is not statistically significant (the standard error against which each bias estimate must be compared is smaller than the quoted standard deviations by a factor of νο). but in any event the overall bias of both parameters is considerably smaller than the associated. random errors. and. will therefore have little effect when added in quadrature.," In all cases, the bias is not statistically significant (the standard error against which each bias estimate must be compared is smaller than the quoted standard deviations by a factor of $\sqrt{6}$ ), but in any event the overall bias of both parameters is considerably smaller than the associated random errors, and will therefore have little effect when added in quadrature."
+ These errors are still larger than the estimated requirements for future surveys (1.8. a~2—4«102. as described in refsec:intro)).," These errors are still larger than the estimated requirements for future surveys (i.e. $\sigma \sim 2-4\times10^{-3}$, as described in \\ref{sec:intro}) )."
+ For cross-ccorrelation techniques to meet these requirements. this excess error will need to be reduced.," For correlation techniques to meet these requirements, this excess error will need to be reduced."
+ We discuss a few options for this 1n $4.1., We discuss a few options for this in $\S\ref{sec:correrrors}$.
+" A number of choices we have made on how to model and measure correlation function parameters (e.g. using a fit for the dependence of the spectroscopic sample's autocorrelation parameters on z vs. using the values for a given z-bin directly: assuming 7p,»erg, VS. a constant 7p: or allowing 7,50) to decrease with redshift vs. forcing a constant 7,,) can affect"," A number of choices we have made on how to model and measure correlation function parameters (e.g. using a fit for the dependence of the spectroscopic sample's autocorrelation parameters on z vs. using the values for a given z-bin directly; assuming $r_{0,pp} \propto r_{0,ss}$ vs. a constant $r_{0,pp}$; or allowing $\gamma_{sp}(z)$ to decrease with redshift vs. forcing a constant $\gamma_{sp}$ ) can affect"
+updated from ? to 40.2XE4.0 d and that of HD 22049 from ? to 11.34+ 1.1d. We also present the rotation period of HD 130322. Poo=26.1$3.5 d. which has not previously been reported.,"updated from \citet{Bal96} to $40.2 \pm 4.0$ d and that of HD 22049 from \citet{Don96} to $11.3 \pm 1.1$ d. We also present the rotation period of HD 130322, $P_{rot}=26.1 \pm 3.5$ d, which has not previously been reported."
+ The seven other stars show periodicities in only one or two seusons so the rotation period is less firmly determined (HD 9826. HD 10697. HD 6930. HD 39744. HD 92788. HD 154345 and HD 217014).," The seven other stars show periodicities in only one or two seasons so the rotation period is less firmly determined (HD 9826, HD 10697, HD 6930, HD 89744, HD 92788, HD 154345 and HD 217014)."
+ Four of these stars” rotation periods have not previously been reported CHD10697. HD 69830. HD 92788 and HD 154345). while the grade assigned to the rotation period of HD 217014 ais been raised from to (2)..," Four of these stars' rotation periods have not previously been reported (HD10697, HD 69830, HD 92788 and HD 154345), while the grade assigned to the rotation period of HD 217014 has been raised from to \citep{Henry00}."
+ Further monitoring and analysis of the H K or photometric fluxes of these stars could increase the reliability of the estimated rotation periods., Further monitoring and analysis of the H K or photometric fluxes of these stars could increase the reliability of the estimated rotation periods.
+ one of the rotation periods refute the inference of an orbiting companion in favourof rotational modulation., None of the rotation periods refute the inference of an orbiting companion in favour of rotational modulation.
+ The inclination of the stellar rotation axes was calculated. qowever the planets presented here are non-transiting. therefore other methods must be used to determine their orbital inclination.," The inclination of the stellar rotation axes was calculated, however the planets presented here are non-transiting, therefore other methods must be used to determine their orbital inclination."
+" It is hoped that future astrometric observations will be able to determine and compare /,, with the inclination of the stellar rotation axes reported here and to infer the line-of-sight alignment of these svstems in order to constrain theories of planet formation and evolution.", It is hoped that future astrometric observations will be able to determine and compare $i_{p}$ with the inclination of the stellar rotation axes reported here and to infer the line-of-sight alignment of these systems in order to constrain theories of planet formation and evolution.
+ Planetary masses were inferred under the assumption that the stellar rotation and planetary orbital axes are aligned., Planetary masses were inferred under the assumption that the stellar rotation and planetary orbital axes are aligned.
+ This method allowed the maximum mass of the orbiting companions to be constrained for most of the systems., This method allowed the maximum mass of the orbiting companions to be constrained for most of the systems.
+ All but one of the companions have a calculated mass below the brown dwarf limit. thereby supporting the inference of their planetary nature.," All but one of the companions have a calculated mass below the brown dwarf limit, thereby supporting the inference of their planetary nature."
+ The calculated mass of HD 92788 b. 28 M;. suggests that it may be a low-mass brown dwarf and warrants further investigation.," The calculated mass of HD 92788 b, 28 $M_{J}$, suggests that it may be a low-mass brown dwarf and warrants further investigation."
+" 500, +100, plus the standard Gaussian case /πτ,=0.","$\pm 500$ , $\pm 100$, plus the standard Gaussian case $f_{\rm NL}=0$."
+ Notice that in this work we prefer to make use of this set of N-body simulations in spite of the more recent one presented in ⋅⋅, Notice that in this work we prefer to make use of this set of N-body simulations in spite of the more recent one presented in \cite{Grossi2009.1}.
+ The reason for this choice is twofold., The reason for this choice is twofold.
+" First, the mass resolution is better and this allows us to have more robust results on small scales, in particular for shear and flexions; second, the assumed values for the parameter fwr span a larger range, covering also cases where the weak lensing signal produced by primordial non-Gaussianity is larger."," First, the mass resolution is better and this allows us to have more robust results on small scales, in particular for shear and flexions; second, the assumed values for the parameter $f_{\rm NL}$ span a larger range, covering also cases where the weak lensing signal produced by primordial non-Gaussianity is larger."
+" To construct the mock light-cones used to perform the ray-tracing simulations, we follow the same procedure described in [Pace (2007), to which we refer for more details."," To construct the mock light-cones used to perform the ray-tracing simulations, we follow the same procedure described in \cite{Pace2007.1}, to which we refer for more details."
+" Since the differentetal] outputs of an N-body simulation represent the redshift evolution of the same initial matter distribution, we need to apply a specific procedure to avoid the introduction of biases related to the fact that the same structures appear approximately at the same positions in different outputs."," Since the different outputs of an N-body simulation represent the redshift evolution of the same initial matter distribution, we need to apply a specific procedure to avoid the introduction of biases related to the fact that the same structures appear approximately at the same positions in different outputs."
+ For this reason we randomly shift and rotate the particle positions exploiting the periodicity of boundaries., For this reason we randomly shift and rotate the particle positions exploiting the periodicity of boundaries.
+" As said before, given our choice for the output redshifts, they overlap for 50 per cent of their comoving side-length, therefore we can just consider particles in the lower half of the rotated and translated boxes."," As said before, given our choice for the output redshifts, they overlap for $50$ per cent of their comoving side-length, therefore we can just consider particles in the lower half of the rotated and translated boxes."
+" Once selected the particles to be used, we project them perpendicularly to the line-of-sight on a regular two-dimensional grid and compute the projected mass density field using the triangular-shaped-cloud (TSC) mass assignment (Hockney&East-wood! 1988)."," Once selected the particles to be used, we project them perpendicularly to the line-of-sight on a regular two-dimensional grid and compute the projected mass density field using the triangular-shaped-cloud (TSC) mass assignment \citep{Hockney1988}."
+. Finally via FFT techniques it is possible to recover the lensing potential associated with the considered matter distribution., Finally via FFT techniques it is possible to recover the lensing potential associated with the considered matter distribution.
+ More details on the numerical procedure will be given in the next section., More details on the numerical procedure will be given in the next section.
+" Our mock light-cones extend along the line-of-sight up to redshift z;=4, using Nout=20 outputs."," Our mock light-cones extend along the line-of-sight up to redshift $z_s=4$, using $N_{\rm out}=20$ outputs."
+" As all the distances are in comoving units, the opening angle of the ray-tracing simulation can be calculated using the comoving distance of the last plane in the stack, resulting in ϐ—5.876 degrees."," As all the distances are in comoving units, the opening angle of the ray-tracing simulation can be calculated using the comoving distance of the last plane in the stack, resulting in $\theta=5.876$ degrees."
+" Using a grid of 2048? points, the corresponding angular resolution of the produced maps is 10.33 arcsec."," Using a grid of $2048^2$ points, the corresponding angular resolution of the produced maps is 10.33 arcsec."
+" Ray-tracing simulations consist in tracing back a bundle of light rays through the matter distribution of the light-cone, from the observer to the sources, which we place at redshift z,=4."," Ray-tracing simulations consist in tracing back a bundle of light rays through the matter distribution of the light-cone, from the observer to the sources, which we place at redshift $z_s=4$."
+" Projected mass maps are converted to projected contrast maps 6P™°}+ by where A; is the the area of the grid cell on the i-th plane, Mi, is the mass projected on the grid cell with indices (1,m) belonging to the i-th plane and [η=250 Mpc/h is the depth of the 2-th subvolume used to build mock cone."," Projected mass maps are converted to projected density-contrast maps $\delta^{{\rm proj,i}}_{\rm lm}$ by where $A_i$ is the the area of the grid cell on the $i$ -th plane, $M^{\rm i}_{\rm lm}$ is the mass projected on the grid cell with indices $(l,m)$ belonging to the $i$ -th plane and $L_{\rm i}=250$ Mpc/h is the depth of the $i$ -th subvolume used to build mock cone."
+ The density contrast is deliberately defined such as to have the unit of a length (see, The density contrast is deliberately defined such as to have the unit of a length \citep[see][]{Hamana2001}.
+ In the previous equation p represents the average& comoving [HamanaMellier/2001)..density of the Universe., In the previous equation $\bar\rho$ represents the average comoving density of the Universe.
+" The lensing potential on each plane W; is related to the projected density contrast through the two-dimensional Poisson equation, namely: Exploiting the periodic boundary conditions of the projected maps, Eq. (9))"," The lensing potential on each plane $\Psi_i$ is related to the projected density contrast through the two-dimensional Poisson equation, namely: Exploiting the periodic boundary conditions of the projected maps, Eq. \ref{eq:poisson}) )"
+ can be solved adopting the FFT techniques., can be solved adopting the FFT techniques.
+" Having the potential on each plane, the lensing quantities can be derived adopting standard finite-difference schemes."," Having the potential on each plane, the lensing quantities can be derived adopting standard finite-difference schemes."
+" In order to perform the ray-tracing simulations, we need to apply the multiple-plane theory to compute the total effect taking into account the contributions from each single lensing plane."," In order to perform the ray-tracing simulations, we need to apply the multiple-plane theory to compute the total effect taking into account the contributions from each single lensing plane."
+ A light-ray is deflected on each plane by the amount d;(61). thus the total deflection is given by the sum of all contributions.," A light-ray is deflected on each plane by the amount $\vec\alpha_i(\vec\theta_i)$, thus the total deflection is given by the sum of all contributions."
+" In particular, if the light-cone is sampled into Nout lens planes and the sources are located on the Nout+1 plane, the relation giving the deflection angle on the 7-th plane of a ray with image position θι reads (in comoving units) here W,(x) is the unscaled lensing potential, ie. the Newtonian potential projected along the line-of-sight, fx is a function depending on the cosmology, w is the comoving distance and a represents the scale factor of the lens plane."," In particular, if the light-cone is sampled into $N_{\rm out}$ lens planes and the sources are located on the $N_{\rm out}+1$ plane, the relation giving the deflection angle on the $i$ -th plane of a ray with image position $\vec{\theta}_1$ reads (in comoving units) here $\Psi_k(\vec{x})$ is the unscaled lensing potential, i.e. the Newtonian potential projected along the line-of-sight, $f_K$ is a function depending on the cosmology, $w$ is the comoving distance and $a$ represents the scale factor of the lens plane."
+" Note that in general the light-rays will intercept the lens planes at arbitrary points, while the potential is defined only at the grid points."," Note that in general the light-rays will intercept the lens planes at arbitrary points, while the potential is defined only at the grid points."
+" Thus, it is necessary to use a bi-linear interpolation to compute the lensing quantities."," Thus, it is necessary to use a bi-linear interpolation to compute the lensing quantities."
+ Differentiating Eq. (10) , Differentiating Eq. \ref{eq:defangle}) )
+"with respect to δι and defining A;=00; and U; the matrix( containing the second derivatives of the lensing potential, one obtains In the previous equation, J represents the identity matrix."," with respect to $\vec{\theta}_1$ and defining $A_i\equiv \partial\vec{\theta}_i/\partial\vec{\theta}_1$ and $U_k$ the matrix containing the second derivatives of the lensing potential, one obtains In the previous equation, $I$ represents the identity matrix."
+" On the source plane, the Jacobian matrix An,,,,+1 is given by where « is now the effective convergence and y=γι+772 is the effective shear."," On the source plane, the Jacobian matrix $A_{N_{\rm out}+1}$ is given by where $\kappa$ is now the effective convergence and $\gamma=\gamma_1+i\gamma_2$ is the effective shear."
+" The term w, called rotation, representsthe asymmetry introduced by multiple lenses."," The term $\omega$, called rotation, representsthe asymmetry introduced by multiple lenses."
+ Differentiating Eq. , Differentiating Eq. \ref{eq:jac}) )
+"with respect to 6;, a recursive relation for the two flexions can be (11))obtained: where Gu=VU is a tensor containing the third derivatives of the lensing potential."," with respect to $\vec{\theta}_i$, a recursive relation for the two flexions can be obtained: where $G_U=\nabla_{\vec{x}}U$ is a tensor containing the third derivatives of the lensing potential."
+" On the source plane, the tensorD reads as where, as for the Jacobian matrix, w; and w» represent the asymmetric terms."," On the source plane, the tensor$D$ reads as where, as for the Jacobian matrix, $\omega_1$ and $\omega_2$ represent the asymmetric terms."
+ In Eq. , In Eq. \ref{eq:Dmatrix}) )
+"we do not take into account the additional terms, called twist (14)and turn, introduced by but several tests performed at different resolutions assure us that (2009),,their inclusion would not significantly affect our results."," we do not take into account the additional terms, called twist and turn, introduced by \cite{Bacon2009}, , but several tests performed at different resolutions assure us that their inclusion would not significantly affect our results."
+original X-ray images was also summed. consisting of both source and background counts.,"original X-ray images was also summed, consisting of both source and background counts."
+ A 3-sigma upper limit to the number of source counts was calculated using the Bayesian approach formulated by Kraft. Burrows Nousek (19913.," A 3-sigma upper limit to the number of source counts was calculated using the Bayesian approach formulated by Kraft, Burrows Nousek (1991)."
+ The choice of a 20 aresee radius follows the work of Carrera et al. (, The choice of a 20 arcsec radius follows the work of Carrera et al. (
+2007). who showed that “aperture photometry” using this radius gave count rates that closely matched those found inEDETECT_CHAIN.,"2007), who showed that “aperture photometry” using this radius gave count rates that closely matched those found in."
+ The upper limits to the X-ray count were determined by dividing by the average exposure time within the same circular area., The upper limits to the X-ray count were determined by dividing by the average exposure time within the same circular area.
+ Using this technique we found that where objects were covered by both PN and MOS data. that the PN data were at least a factor of two more sensitive.," Using this technique we found that where objects were covered by both PN and MOS data, that the PN data were at least a factor of two more sensitive."
+ Rather than attempt to combine the results from the three instruments we have either (a) taken the upper limit from the PN. where PN data are present (23 objects). (b) taken the the upper limit from one MOS detector where an object was only covered by that detector (3 objects) or (c) taken the average upper limit from both MOS detectors (2 objects) and divided by v2. as in these cases the two upper limits were very similar.," Rather than attempt to combine the results from the three instruments we have either (a) taken the upper limit from the PN, where PN data are present (23 objects), (b) taken the the upper limit from one MOS detector where an object was only covered by that detector (3 objects) or (c) taken the average upper limit from both MOS detectors (2 objects) and divided by $\sqrt{2}$, as in these cases the two upper limits were very similar."
+" The count rate upper limits were converted to upper limits in X-ray flux and upper limits to ἐνLi,,; using the procedures deseribed in the previous subsection.", The count rate upper limits were converted to upper limits in X-ray flux and upper limits to $L_x/L_{\rm bol}$ using the procedures described in the previous subsection.
+ The details of the upper limit measurements are presented separately in Table 2 (available in electronic form only)., The details of the upper limit measurements are presented separately in Table 2 (available in electronic form only).
+" The uncertainties in the logLa/Li,,, values quoted in Table | incorporate the statistical count rate uncertainties and. the systematic instrument response uncertainties mentioned in section 2.3.", The uncertainties in the $\log L_{\rm x}/L_{\rm bol}$ values quoted in Table 1 incorporate the statistical count rate uncertainties and the systematic instrument response uncertainties mentioned in section 2.3.
+ In addition we have included (in quadrature) uncertainties in. the bolometric flux due to the photometric uncertainties (or variability) implied by equations | and 2. which amount to about 0.07 dex.," In addition we have included (in quadrature) uncertainties in the bolometric flux due to the photometric uncertainties (or variability) implied by equations 1 and 2, which amount to about $\pm
+0.07$ dex."
+ The tinal quoted uncertainties range from 0.08 dex to 0.21 dex. with a mean of 0.10 dex.," The final quoted uncertainties range from 0.08 dex to 0.21 dex, with a mean of 0.10 dex."
+ Other contributing uncertainties have also been considered., Other contributing uncertainties have also been considered.
+ The X-ray flux conversion factors are rather insensitive to the details ofthe spectral model., The X-ray flux conversion factors are rather insensitive to the details of the spectral model.
+ For instance doubling the temperature of the hot component increases the conversion factor by just 3 per cent and the HR increases to -0.39: varying the metal abundance in the range 0.1]«Z<1.0 changes the conversion factor by only £5 per cent: doubling the ratio of hot-to-cool plasma emission measures increases the conversion factor by 3 per cent and increases the HR to -0.32., For instance doubling the temperature of the hot component increases the conversion factor by just 3 per cent and the HR increases to -0.39; varying the metal abundance in the range $0.1<Z<1.0$ changes the conversion factor by only $\pm 5$ per cent; doubling the ratio of hot-to-cool plasma emission measures increases the conversion factor by 3 per cent and increases the HR to -0.32.
+ The value assumed for Αμ has a little more effect., The value assumed for $N_H$ has a little more effect.
+ The value of Ng—3.107 7 is derived from the cluster reddening of (D.V)=0.06+0.02 (Clariá 1982) and the relationship between reddening and Αμ found by Bohlin. Savage Drake (1978). and could be uncertain by factors of two.," The value of $N_H =
+3\times10^{20}$ $^{-2}$ is derived from the cluster reddening of $E(B-V)=0.06\pm 0.02$ (Clariá 1982) and the relationship between reddening and $N_H$ found by Bohlin, Savage Drake (1978), and could be uncertain by factors of two."
+" Altering Vig from our assumed value to either 1.5107 7 or 6O7"" > would lead to conversion factors about 5 per cent smaller or 10 per cent larger respectively and HR would change to either -O. or -0.35 respectively.", Altering $N_H$ from our assumed value to either $1.5\times 10^{20}$ $^{-2}$ or $6\times 10^{20}$ $^{-2}$ would lead to conversion factors about 5 per cent smaller or 10 per cent larger respectively and HR would change to either -0.44 or -0.35 respectively.
+ All these uncertainties are small compared with those we have already considered and they are neglected., All these uncertainties are small compared with those we have already considered and they are neglected.
+ The final source of uncertainty is difficult to quantify. but is probably dominant when considering a single epoch of X-ray data. namely the coronal variability of active stars.," The final source of uncertainty is difficult to quantify, but is probably dominant when considering a single epoch of X-ray data, namely the coronal variability of active stars."
+ Active. low-mass stars show frequent X-ray flaring behaviour on timescales of minutes and hours. resulting in an upward bias in an X-ray flux estimated. as here. over the course of more than a day.," Active, low-mass stars show frequent X-ray flaring behaviour on timescales of minutes and hours, resulting in an upward bias in an X-ray flux estimated, as here, over the course of more than a day."
+ The more active the low-mass star. the more frequently it exhibits large ray flares CAudard et al.," The more active the low-mass star, the more frequently it exhibits large X-ray flares (Audard et al."
+ 2000)., 2000).
+ Rather than try to correct for this. and as the comparison samples in section + have not had any flare exclusion applied. we continue to use the time-averaged X-ray flux to represent coronal activity.," Rather than try to correct for this, and as the comparison samples in section 4 have not had any flare exclusion applied, we continue to use the time-averaged X-ray flux to represent coronal activity."
+ An idea of the uncertainties can be gained by looking at similar estimates from more than one epoch., An idea of the uncertainties can be gained by looking at similar estimates from more than one epoch.
+ For example. Gagné.. Caillault Stauffer (1995) tind that young Pleiades low-mass stars show differences of more than a factor of two in their X-ray fluxes only 25 per cent of the time on timescales of I-10 years.," For example, Gagné,, Caillault Stauffer (1995) find that young Pleiades low-mass stars show differences of more than a factor of two in their X-ray fluxes only 25 per cent of the time on timescales of 1–10 years."
+ Marino et al. (, Marino et al. (
+2003) find that the median level of X-ray flux variability of. K3-M dwarfs in the Pleiades is abou ddex on timescales of months.,2003) find that the median level of X-ray flux variability of K3-M dwarfs in the Pleiades is about dex on timescales of months.
+ Jeffries et al. (, Jeffries et al. (
+2006) looked a the variability of G- to M-dwarfs in NGC 2547 itselfon timescales of 7 years. finding a median absolute deviation from equal X-ray luminosity of about ddex in G- and K dwarfs. with a hint tha the M-dwarfs are slightly more variable.,"2006) looked at the variability of G- to M-dwarfs in NGC 2547 itself on timescales of 7 years, finding a median absolute deviation from equal X-ray luminosity of about dex in G- and K dwarfs, with a hint that the M-dwarfs are slightly more variable."
+ Only about 20 per cent of sources varied by more than a factor of two., Only about 20 per cent of sources varied by more than a factor of two.
+ Our conclusion is tha we should assume an additional uncertainty of about ddex in our single-epoch measurements. but caution that the distribution is probably non-Gaussian in the sense that a sample will probably contain a small number of objects that are upwardly biased by more than a factor of two by large flares.," Our conclusion is that we should assume an additional uncertainty of about dex in our single-epoch measurements, but caution that the distribution is probably non-Gaussian in the sense that a sample will probably contain a small number of objects that are upwardly biased by more than a factor of two by large flares."
+ In Figure | we show the dependence of the HR. measured by the PN camera. on the intrinsic colour of the star and on the magnetic activity level. expressed as {ήLia.," In Figure 1 we show the dependence of the HR, measured by the PN camera, on the intrinsic colour of the star and on the magnetic activity level, expressed as $L_x/L_{\rm bol}$."
+ There is no significant dependence of HR on colour. and hence the approximation of a uniform flux conversion factor with spectral type is shown to be reasonable.," There is no significant dependence of HR on colour, and hence the approximation of a uniform flux conversion factor with spectral type is shown to be reasonable."
+ The weighted mean HR is —.0.42+0.02. with a standard deviation of 0.16.," The weighted mean HR is $-0.42\pm 0.02$, with a standard deviation of 0.16."
+ The standard deviation is about twice that expected from the statistical errors. hinting at some genuine HR variation.," The standard deviation is about twice that expected from the statistical errors, hinting at some genuine HR variation."
+ Note that this assumes that the HR. uncertainties are symmetric and Gaussian. even though the X-ray data are Poissonian in nature and HR is strictly limited to 1 <HHR<1 (e.g. see Park et al.," Note that this assumes that the HR uncertainties are symmetric and Gaussian, even though the X-ray data are Poissonian in nature and HR is strictly limited to $-1 \leq$ $\leq 1$ (e.g. see Park et al."
+ 2006)., 2006).
+ The majority of sources are detected with sufficiently large numbers of counts in both energy ranges and with sufficiently precise count rates to make such an assumption reasonable., The majority of sources are detected with sufficiently large numbers of counts in both energy ranges and with sufficiently precise count rates to make such an assumption reasonable.
+ The second plot shows there is some evidence that HR increases slightly with activity level., The second plot shows there is some evidence that HR increases slightly with activity level.
+ A best-fit linear model is HR=0.00(-0.12).|O.11(zE0.04)logLis. but the residuals suggest there may still be star-to-star variationLf at a given activity level.," A best-fit linear model is $=-0.09(\pm 0.12) + 0.11(\pm 0.04) \log L_x/L_{\rm bol}$, but the residuals suggest there may still be star-to-star variation at a given activity level."
+ Note though that time variability of the X-ray activity or HR values has not been considered (see section 2.5) and this might plausibly account for some of this star-to-star variation., Note though that time variability of the X-ray activity or HR values has not been considered (see section 2.5) and this might plausibly account for some of this star-to-star variation.
+ A significant increase in average coronal temperatures cand ence HR) with activity level is well established in field stars (e.g. Telleschi et al., A significant increase in average coronal temperatures (and hence HR) with activity level is well established in field stars (e.g. Telleschi et al.
+ 2005)., 2005).
+ The relationship between HR and LjLi ound here and the size of the intrinsic scatter are indistinguishable rom those found for somewhat higher mass G- and K-stars in GC 2547 by Jeffries et al. (, The relationship between HR and $L_x/L_{\rm bol}$ found here and the size of the intrinsic scatter are indistinguishable from those found for somewhat higher mass G- and K-stars in NGC 2547 by Jeffries et al. (
+2006). as is the likelihood of an intrinsic scatter.,"2006), as is the likelihood of an intrinsic scatter."
+ The mean HR agrees well (by design) with the HR predicted by the spectral model used to calculate the X-ray Πιιχος., The mean HR agrees well (by design) with the HR predicted by the spectral model used to calculate the X-ray fluxes.
+ The probable intrinsic variation in AR ata given ἐνLia corresponds to variations of about a factor of two in the emission measure ratio of the hot and cool plasmas in the two-temperature coronal model., The probable intrinsic variation in HR at a given $L_{x}/L_{\rm bol}$ corresponds to variations of about a factor of two in the emission measure ratio of the hot and cool plasmas in the two-temperature coronal model.
+ Such variations lead to very small uncertainties in LfLi of order a few per cent (see section 2.5) and so we have not attempted to correct for them., Such variations lead to very small uncertainties in $L_{x}/L_{\rm bol}$ of order a few per cent (see section 2.5) and so we have not attempted to correct for them.
+We adopted model accretion disks lrom the optically thick disk model exid of (1998).,We adopted model accretion disks from the optically thick disk model grid of \citet{Wade1998}.
+". In these accretion disk models. the innermost disk radius. Hj. is fixed at a fractional white dwail radius of a=R;/HR,4;-1.05."," In these accretion disk models, the innermost disk radius, $_{in}$, is fixed at a fractional white dwarf radius of $x = R_{in}/R_{wd} = 1.05$."
+" The outermost disk radius. ,,;. was chosen so that ντεο) is near 100001 since disk annuli bevond this point. which are cooler zones with larger radii. would provide only a very small contribution to the mid ancl far UV disk flux. particularly the SWP FUV bandpass."," The outermost disk radius, $_{out}$ , was chosen so that $_{eff}(R_{out})$ is near 10,000K since disk annuli beyond this point, which are cooler zones with larger radii, would provide only a very small contribution to the mid and far UV disk flux, particularly the SWP FUV bandpass."
+ The mass transfer rate is assumed to be the same for all radii., The mass transfer rate is assumed to be the same for all radii.
+" Thus. (he run of disk temperature with radius is taken to be: where 7—r/I,,; and oT!=3GM,4MΣπ, The disk is divided into a set of ring annuli."," Thus, the run of disk temperature with radius is taken to be: where $x = r/R_{wd}$ and $\sigma T_{s}^{4} = 3 G M_{wd}\dot{M}/8\pi R_{wd}^{3}$ The disk is divided into a set of ring annuli."
+ The vertical structure of each ring is computed with TLUSDISIx (IIubenyv1990).. which is a derivative of the stellar atmosphere program TLUSTY (IlubenvLoss).," The vertical structure of each ring is computed with TLUSDISK \citep{Hubeny1990}, which is a derivative of the stellar atmosphere program TLUSTY \citep{Hubeny1988}."
+.. The spectrum synthesis program SYNSPEC described bv Hubeny&Lanz(1995) is used to solve the radiative transfer equation to compute the local. rest. [rame spectrum for each ring of the disk.," The spectrum synthesis program SYNSPEC described by \citet{Hubeny1995} is used to solve the radiative transfer equation to compute the local, rest frame spectrum for each ring of the disk."
+ In addition to detailed profiles of the ll and Ile lines. (he spectrum synthesis includes metal lines up to Nickel (Z = 28).," In addition to detailed profiles of the H and He lines, the spectrum synthesis includes metal lines up to Nickel (Z = 28)."
+ The accretion disks are computed in LTE and the chemical composition of the accretion disk is kept fixed at solar values in our study., The accretion disks are computed in LTE and the chemical composition of the accretion disk is kept fixed at solar values in our study.
+ Limb darkening of the disk is fully taken into account in the manner described by Diaz involving the Eddineton-Barbier relation. the increase of kinetic temperature with depth in the disk. aud the wavelength. and temperature dependence of the Planck [unction.," Limb darkening of the disk is fully taken into account in the manner described by \citet{Diaz1996} involving the Eddington-Barbier relation, the increase of kinetic temperature with depth in the disk, and the wavelength and temperature dependence of the Planck function."
+ The details of our 472 (per degree of freedom) minimization fitting procedures are discussed in detail in Sionetal.(1995)., The details of our $\chi^{2}_{\nu}$ (per degree of freedom) minimization fitting procedures are discussed in detail in \citet{Sion1995}.
+. We (take anv reliable published parameters like ihe WD mass. or orbital inclination onlv as an initial guess in searching for the best-fitting accretion disk models.," We take any reliable published parameters like the WD mass, or orbital inclination only as an initial guess in searching for the best-fitting accretion disk models."
+ We relax all constraints except the lower limit Ixnigge(2006) distance., We relax all constraints except the lower limit \citet{Knigge2006} distance.
+ For each systemss spectrum. we carry oul [its for every combination of M. inclination and white dwarf mass in the Wade&IIubeny.(1998) library.," For each systems's spectrum, we carry out fits for every combination of $\dot{M}$, inclination and white dwarf mass in the \citet{Wade1998} library."
+ The values of are 15. 41. 60. 75 and SI.," The values of are 18, 41, 60, 75 and $81\degr$."
+ The range of accretion rates covers —10.5<logM.—8.0 in steps of 0.5 in the log and five different. values of the white dwarf mass. namely. 0.4. 0.55. 0.3. 1.0. and 1.2 M...," The range of accretion rates covers $-10.5 < \log \dot{M} < -8.0$ in steps of 0.5 in the log and five different values of the white dwarf mass, namely, 0.4, 0.55, 0.8, 1.0, and 1.2 $_{\sun}$."
+ Out of the roughly 900 models using every combination of 7. M and Mj. we try to isolate the accretion disk model.," Out of the roughly 900 models using every combination of $i$, Ṁ and $_{wd}$, we try to isolate the best-fitting accretion disk model."
+ In effect. we search for the best fit based upon(1) the minimum V value achieved. (2) the goodness of fit of the continuum slope. (3) the goodness of fit to the," In effect, we search for the best fit based upon(1) the minimum $\chi^{2}$ value achieved, (2) the goodness of fit of the continuum slope, (3) the goodness of fit to the"
+For isotopes with no primordial composition JV?=0.,For isotopes with no primordial composition $X_i^o=0$.
+ After some simple algebra from equations (2) and (3). the chemical enrichment can be written.," After some simple algebra from equations (2) and (3), the chemical enrichment can be written."
+ where y; is the stellar vield defined as the mass fraction of the total mass returned by a generation of stars in the form of (he isotope 7., where $y_i$ is the stellar yield defined as the mass fraction of the total mass returned by a generation of stars in the form of the isotope $i$.
+ The results from the combined effect of the stellar vields and IMFESs for Population II] are summarized in Table 1., The results from the combined effect of the stellar yields and IMFs for Population III are summarized in Table 1.
+ The second column gives the total metallicity Z reached in (he material ejected by the star burst (i.e. without considering the dilution with the metal-[ree IGM gas) relative to the solar metallicity (7.= 0.0139).," The second column gives the total metallicity $Z$ reached in the material ejected by the star burst (i.e., without considering the dilution with the metal-free IGM gas) relative to the solar metallicity $Z_\odot=0.0189$ )."
+ Column three shows (he maximum efficient factor p allowable for each IAIF so as not to surpass the maximum metallicity observed in the Ly-a [orest. Z~4x107Z..," Column three shows the maximum efficient factor $p$ allowable for each IMF so as not to surpass the maximum metallicity observed in the $\alpha$ forest, $Z\sim 4\times 10^{-3} Z_\odot$."
+ Note that quite low values of p are allowed lor all the IMESs., Note that quite low values of $p$ are allowed for all the IMFs.
+ The most restrictive p values are those obtained for the NUa NUb IMEs because of the high metal pollution from VMOs (Z/Z.>1. see column(νο)1.," The most restrictive $p$ values are those obtained for the NUa NUb IMFs because of the high metal pollution from VMOs $Z/Z_\odot>1$, see column."
+ These low values of p contrast wilh the current estimate of the efficient [actor in regions of star formation in the Galaxy. p<107. or in other spiral galaxies (Nennicutt|1993).," These low values of $p$ contrast with the current estimate of the efficient factor in regions of star formation in the Galaxy, $p\la 10^{-2}$, or in other spiral galaxies \citep{ken98}."
+. llowever. for anv IME choice in Table 1 the corresponding p [actor is large enough for the," However, for any IMF choice in Table 1 the corresponding $p$ factor is large enough for the"
+(Sect. 3.2)),(Sect. \ref{hostnuc}) )
+ leads to a correlation between black hole mass and nuclear luminosity: however. when we include the error bars in the black hole masses. the correlation is not statistically significant.," leads to a correlation between black hole mass and nuclear luminosity; however, when we include the error bars in the black hole masses, the correlation is not statistically significant."
+" Given AGN black hole masses estimated in this way. we can derive the Eddington luminosity: where MM, is the black hole mass. my the mass of the proton. and op the Thompson "," Given AGN black hole masses estimated in this way, we can derive the Eddington luminosity: where $M_{\bullet}$ is the black hole mass, $m_p$ the mass of the proton, and $\sigma_T$ the Thompson cross-section."
+From this we ealeulate the Eddington ratio using the corrected nuclear magnitudes. assumine a spectral index ofa=1 to estimate bolometric Iuminositv.," From this we calculate the Eddington ratio using the corrected nuclear magnitudes, assuming a spectral index of $\alpha=1$ to estimate bolometric luminosity."
+ The nuclear magnitudes of DL Lac objects are corrected. [or beaming using lower limits to the Doppler factors. meaning the calculated IExldington ratios are upper limits.," The nuclear magnitudes of BL Lac objects are corrected for beaming using lower limits to the Doppler factors, meaning the calculated Eddington ratios are upper limits."
+ For the eight BL Lacs with measured Doppler [actor limits. the median Eddington ratio limit is naCOS<0.002.," For the eight BL Lacs with measured Doppler factor limits, the median Eddington ratio limit is $\frac{L_{bol.}}{L_{Edd}} \lesssim
+0.002$."
+ For the entire low-power sample. corrected with the median Doppler factor. we obtain =~Leas<3x10!.," For the entire low-power sample, corrected with the median Doppler factor, we obtain $\frac{L_{bol.}}{L_{Edd}} \lesssim 3\times10^{-4}$."
+ The median Eddingtono ratio for the high-powero sources is much higher. with LewLoot.—(0.1.," The median Eddington ratio for the high-power sources is much higher, with $\frac{L_{bol.}}{L_{Edd}} = 0.1$."
+ The upper part of Figure 5. shows the histogram of Edclington ratios (or Ecldineton ratio limits) lor the reduced low-power sample ancl the hieh-power sample., The upper part of Figure \ref{fig5} shows the histogram of Eddington ratios (or Eddington ratio limits) for the reduced low-power sample and the high-power sample.
+" Black hole mass appears to be much more tightly correlated wilh velocity dispersion (o,) than with bulge luminosity (Gebhardtefa£.2000:Ferrarese&Gebhardt 2001)."," Black hole mass appears to be much more tightly correlated with velocity dispersion $\sigma_e$ ) than with bulge luminosity \citep{Gebhardt, Ferrarese, KormendyG}."
+".. If we assume that our host galaxies lie on the Fundamental Plane lor normal ellipticals. we can infer velocity dispersions using the measured values of 7. ancl (1), the mean surface brightness within r,."," If we assume that our host galaxies lie on the Fundamental Plane for normal ellipticals, we can infer velocity dispersions using the measured values of $r_e$ and $\langle I \rangle_e$ — the mean surface brightness within $r_e$."
+ This assumption is reasonable. as we know that these host galaxies are indistinguishable from normal elliptical galaxies morphologically. aud from 3.1 we know that they are consistent dvnamically. with their Kormendy relations matching.," This assumption is reasonable, as we know that these host galaxies are indistinguishable from normal elliptical galaxies morphologically, and from \ref{host} we know that they are consistent dynamically, with their Kormendy relations matching."
+ We use (he Fundamental Plane parameters measured by from R-band photometry of a Large sample of E ancl SO galaxies across several clusters:," We use the Fundamental Plane parameters measured by \citet{Jorgensen}
+ from R-band photometry of a large sample of E and S0 galaxies across several clusters:"
+it up in the process.,it up in the process.
+ Since the core only has to accrete a small fraction of the angular momentum available from the secondary to be spun up to critical rotation. we may expect that this itself will alfect the eivnanmies of the interaction.," Since the core only has to accrete a small fraction of the angular momentum available from the secondary to be spun up to critical rotation, we may expect that this itself will affect the dynamics of the interaction."
+ To demonstrate this. we performed four simulations for streams with the same initial parameters and similar core properties except that we varied. the rotation of the core from 0 to wwithin a core radius of 4.101em?..," To demonstrate this, we performed four simulations for streams with the same initial parameters and similar core properties except that we varied the rotation of the core from 0 to within a core radius of $4\times 10^{10}\,$."
+" ""Ehe results of these simulations are shown in Figures 5. and 6.."," The results of these simulations are shown in Figures \ref{fig:core}
+ and \ref{fig:core2}. ."
+ s Figure 5 shows. the core.stream interaction changes qualitatively ancl quantitatively. where the penetration depth decreases with increasing core rotation.," As Figure \ref{fig:core} shows, the core–stream interaction changes qualitatively and quantitatively, where the penetration depth decreases with increasing core rotation."
+ In the case of arapiclly rotating core. the penetration depth is centrifusallv limited (i.c. by the angular momentum in the stream).," In the case of a rapidly rotating core, the penetration depth is centrifugally limited (i.e. by the angular momentum in the stream)."
+ In Figure 6 we show the logarithm of the acliabatic constant and the specifie angular momentum near the center of the streams in the four simulations as a function of radius., In Figure \ref{fig:core2} we show the logarithm of the adiabatic constant and the specific angular momentum near the center of the streams in the four simulations as a function of radius.
+ In the two simulations with no or slow core rotation (solid aud dashed. curves). the stream's angular velocity is ellicientlv being braked below a radius 2.6«101 em by its interaction with the core and reaches the deepest. point of penetration (marked as fillecl circles). where the adiabatic constant is comparable to the ambient. one.," In the two simulations with no or slow core rotation (solid and dashed curves), the stream's angular velocity is efficiently being braked below a radius $2.6\times 10^{10}\,$ cm by its interaction with the core and reaches the deepest point of penetration (marked as filled circles) where the adiabatic constant is comparable to the ambient one."
+ The specifie angular monentum is always substantially less than the local critical angular momentum for centrifugal support. given by y“GR. where g is the gravitational acceleration at radius 2.," The specific angular momentum is always substantially less than the local critical angular momentum for centrifugal support, given by $\sqrt{gR^3}$, where $g$ is the gravitational acceleration at radius $R$."
+ On he other hand. in the two cases with rapicl core rotation (clot-clasheel ancl dot-dot-dot-dashed curves). the stream is inst accelerated in the azimuthal direction by the interaction with the core. in fact reaching an angular velocity that is arger than the local velocity of the core (even though the initial angular velocity at the core boundary was less).," On the other hand, in the two cases with rapid core rotation (dot-dashed and dot-dot-dot-dashed curves), the stream is first accelerated in the azimuthal direction by the interaction with the core, in fact reaching an angular velocity that is larger than the local velocity of the core (even though the initial angular velocity at the core boundary was less)."
+ At the »oint. of deepest penetration. the specific angular momentum is close to the critical angular momentum. and hence the stream becomes centrifugally supported.," At the point of deepest penetration, the specific angular momentum is close to the critical angular momentum, and hence the stream becomes centrifugally supported."
+ Note also that. in his case. the stream does not bounce but starts to form a ring orbiting the core and. gradually mixing with it.," Note also that, in this case, the stream does not bounce but starts to form a ring orbiting the core and gradually mixing with it."
+ This demonstrates that the main clleet of rapid. core rotation is to the azimuthal braking of the stream., This demonstrates that the main effect of rapid core rotation is to the azimuthal braking of the stream.
+ While this leads to less entropy generation in the stream (see the top panel in Fig. 6)).," While this leads to less entropy generation in the stream (see the top panel in Fig. \ref{fig:core2}) ),"
+ which ordinarily would mean that the stream should be able to penetrate deeper. it makes the angular momentum of the stream itself a kev factor in limiting the penetration depth.," which ordinarily would mean that the stream should be able to penetrate deeper, it makes the angular momentum of the stream itself a key factor in limiting the penetration depth."
+ In the simulation with the fastest core. the stream matter is still substantially clenser than the ambient. matter at the final point shown in the simulations.," In the simulation with the fastest core, the stream matter is still substantially denser than the ambient matter at the final point shown in the simulations."
+ Since it is also rotating faster than the local core material. it is reasonable to expect that the stream will continue to spiral in as it is being braked by the anibicnt mecium and will penetrate somewhat deeper than is shown in the figures. (," Since it is also rotating faster than the local core material, it is reasonable to expect that the stream will continue to spiral in as it is being braked by the ambient medium and will penetrate somewhat deeper than is shown in the figures. ("
+We were unable to follow the subsequent evolution since the stream matter was leaving the domain of the calculation at this point.),We were unable to follow the subsequent evolution since the stream matter was leaving the domain of the calculation at this point.)
+ Even in the case where the core is non-rotating. the specific angular momentum in the stream: increases during a portion of its trajectory (see the solid. curve in the bottom panel between 3 and 3.8102 em).," Even in the case where the core is non-rotating, the specific angular momentum in the stream increases during a portion of its trajectory (see the solid curve in the bottom panel between 3 and $3.8\times
+10^{10}\,$ cm)."
+ This is a direct consequence of the stream splitting into two components in the impact region. a forward and a backward: component (this can be seen more clearly in Figs 1 and 2).," This is a direct consequence of the stream splitting into two components in the impact region, a forward and a backward component (this can be seen more clearly in Figs 1 and 2)."
+ As matter in the forward component is pushed by matter following from. behind. it is being accelerated and gainsangular momentum.," As matter in the forward component is pushed by matter following from behind, it is being accelerated and gainsangular momentum,"
+of CalFUSE (Dixonetal.2007) following the same procedure as in previous works (Lor details see Godonοἱal... (2009))).,of CalFUSE \citep{dix07} following the same procedure as in previous works (for details see \citet{god09}) ).
+ TheFUSE spectrum of MV Lyr (Figure 1) is characterized by broad and deep absorption lines from highly ionized species ancl (he absence of Hydrogen Lyman lines (except lor sharp lines from the ISM)., The spectrum of MV Lyr (Figure 1) is characterized by broad and deep absorption lines from highly ionized species and the absence of Hydrogen Lyman lines (except for sharp lines from the ISM).
+ These absorption lines are (~923A )) η. (1131.6 )). (1062.6 Jy. (~ D. 5&l ). and )).," These absorption lines are $\sim$ ), ), (1131.6 ), (1062.6 ), $\sim$ ), ), (1122.5 ), and )."
+ Such deep and broad absorption lines are often. observed in CVs in high state. even in svstems will a large inclination where the Ixepleriai velocity broadening is expected to produce a rather smooth continuum (e.g. see theFUSE spectrum of V3885 Ser in (2009))).," Such deep and broad absorption lines are often observed in CVs in high state, even in systems with a large inclination where the Keplerian velocity broadening is expected to produce a rather smooth continuum (e.g. see the spectrum of V3885 Sgr in \citet{lin09}) )."
+ These broad and deep absorption lines most likely form in a hot corona above the heated stellar surface. BL or disk.," These broad and deep absorption lines most likely form in a hot corona above the heated stellar surface, BL or disk."
+ Additional transition wavelengths have been marked on the spectrum in Figure 1 and can be associated. with some spectral leatures: however. ihey are not clearly detected.," Additional transition wavelengths have been marked on the spectrum in Figure 1 and can be associated with some spectral features; however, they are not clearly detected."
+ There are a few sharp absorption lines from the interstellar medium such as1020A.1036À. 1048.2 1066.7À. and1145AÀ.," There are a few sharp absorption lines from the interstellar medium such as, 1048.2 , and."
+ The sharp emission lines ave artifacts due to air glow or direct rellection of sunlight inside the telescope (e... such as the ]lines).," The sharp emission lines are artifacts due to air glow or direct reflection of sunlight inside the telescope (e.g., such as the lines)."
+ The modeling of theFUSE spectrum is carried out in steps., The modeling of the spectrum is carried out in steps.
+ First. spherically svinnetric stellar atmosphere structure as well as the vertical structure of disk rings are computed using the code TLUSTY (IIubeny.1958).," First, spherically symmetric stellar atmosphere structure as well as the vertical structure of disk rings are computed using the code TLUSTY \citep{hub88}."
+. For the stellar structure the input parameters are the surface gravity. effective temperature aud (he surface composition of the star.," For the stellar structure the input parameters are the surface gravity, effective temperature and the surface composition of the star."
+ For the composition it is often enough to take Ivdrogen and Helium aud neglect the metals as long, For the composition it is often enough to take Hydrogen and Helium and neglect the metals as long
+et al.,et al.
+ 2007. Scharimer et al.," 2007, Scharmer et al."
+ 2008b)., 2008b).
+ Due to the presence of flow components directed along as well as across the axis of the fikunecut. such field lines get twisted.," Due to the presence of flow components directed along as well as across the axis of the filament, such field lines get twisted."
+ From the observations alone it is not possible to determine how deep the covective down aud upflows reach below the surface., From the observations alone it is not possible to determine how deep the covective down and upflows reach below the surface.
+ Recenth. Spruit et al. (," Recently, Spruit et al. ("
+2010) ave argued for deep couvection.,2010) have argued for deep convection.
+ Them criticisur of the rol-like couvection interpretation is partly based on the wisconception that these rolls need to have a circular cross-section., Their criticism of the roll-like convection interpretation is partly based on the misconception that these rolls need to have a circular cross-section.
+ Of course. they can extend ιο deeper.," Of course, they can extend much deeper."
+ ILowever. the interpretation of convective rolls docs imply hat at some poit below the surface the sirouuding uaenetie field wraps around the couvective fluid again. separating it frou the feld-free convecting eas below the xumnuabra.," However, the interpretation of convective rolls does imply that at some point below the surface the surrounding magnetic field wraps around the convective fluid again, separating it from the field-free convecting gas below the penumbra."
+ This is iu contrast to the fleld-free eap model of Spruit aud Scharimer (2006). cf.," This is in contrast to the field-free gap model of Spruit and Scharmer (2006), cf."
+ Scharimer (2009) in which the field-free eap is open at the bottom., Scharmer (2009) in which the field-free gap is open at the bottom.
+ Numerical sinmulatious (Ποια et al., Numerical simulations (Heinemann et al.
+ 2008. Repel et al.," 2008, Rempel et al."
+ 2009) indicate that the convectiveται features reach deep. but are indeed wrapped by the surrounding ρααυτα] field and do not open iuto the field-free eas surrouudiug the sunspot.," 2009) indicate that the convective features reach deep, but are indeed wrapped by the surrounding prnumbral field and do not open into the field-free gas surrounding the sunspot."
+ This is the main difference between the two Xetures. besides the absence of a magnetic field within he gaps (Spruit and Scharmer 2006). but its presence in the rolls proposed by Zakharov et al. (," This is the main difference between the two pictures, besides the absence of a magnetic field within the gaps (Spruit and Scharmer 2006), but its presence in the rolls proposed by Zakharov et al. ("
+2008).,2008).
+ The present paper has provided further support for the idea that overturning convection is an important process In αλαπιο the peuunbral temperature., The present paper has provided further support for the idea that overturning convection is an important process in maintaining the penumbral temperature.
+ The clear sienal of an upflow alone the central axis of a bright filament has been reported by Franz Schlichemmaicr (2010) aud Tchimote (2010) but downuflows at its sides has still to be seen. however. although indirect evidence has been found by Mairquez et al. (," The clear signal of an upflow along the central axis of a bright filament has been reported by Franz Schlichenmaier (2010) and Ichimoto (2010) but downflows at its sides has still to be seen, however, although indirect evidence has been found by Márrquez et al. ("
+2006) and Sánuchez Ahneida et al. (,2006) and Sánnchez Almeida et al. (
+2007).,2007).
+ This work has been partly supported by the WCU eraut No., This work has been partly supported by the WCU grant No.
+ R31-10016 funded bv the Ixorean Ministry of Education. Science aud Technology.," R31-10016 funded by the Korean Ministry of Education, Science and Technology."
+ Iinode is à Japanese nussion developed aud launched by ISAS/JANA. collaborating with NAOJ as a domestic partucr. NASA and STFC (UK) as international partners.," Hinode is a Japanese mission developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic partner, NASA and STFC (UK) as international partners."
+ Scieutific operation of the Tinode missiou is conducted bv the lliuode science tea organized at ISAS/JANA., Scientific operation of the Hinode mission is conducted by the Hinode science team organized at ISAS/JAXA.
+ Support for the post-launch operation is provided by JANA aud NAOJ (Japan). STFC (U.K.). NASA (U.S.A.). ESA. and NSC (Norway).," Support for the post-launch operation is provided by JAXA and NAOJ (Japan), STFC (U.K.), NASA (U.S.A.), ESA, and NSC (Norway)."
+surveys. and GRBs would occur wherever the star formation akes place.,"surveys, and GRBs would occur wherever the star formation takes place."
+ However the CRBs may also be biased. tracers of star formation. and the current GRB data suggests that an enhancement of the GAB rate at high-z is necessary o account for the observed redshift cüstribution of GRBs (Virgilietal.2011).," However the GRBs may also be biased tracers of star formation, and the current GRB data suggests that an enhancement of the GRB rate at $z$ is necessary to account for the observed redshift distribution of GRBs \citep{Virgili.etal:11}."
+". A GRB rate increase due to lower metallicity at high-z: may not be the singular cause. and other effect such as the evolution of the GRB luminosity ""unction break with redshift might be important."," A GRB rate increase due to lower metallicity at $z$ may not be the singular cause, and other effect such as the evolution of the GRB luminosity function break with redshift might be important."
+ Civen hese complications. further studies are needed in the future o better understand the relationship between the CRB rate and cosmic SERD.," Given these complications, further studies are needed in the future to better understand the relationship between the GRB rate and cosmic SFRD."
+" Based on the results of self-consistent cosmological ivdrodyvnamic simulations. we argued that the current > observations are missing low-mass galaxies with A,10A4. at z>d when accounting for the total p;. and his results in the inconsistency between the observational estimates of SERD and ον."," Based on the results of self-consistent cosmological hydrodynamic simulations, we argued that the current $z$ observations are missing low-mass galaxies with $\Mstar \lesssim 10^9\Msun$ at $z>4$ when accounting for the total $\rhostar$, and this results in the inconsistency between the observational estimates of SFRD and $\rhostar$."
+ But are. there any other »ossibilities to explaün the above inconsistency. or ds il possible that our simulations are incorrect?," But are there any other possibilities to explain the above inconsistency, or is it possible that our simulations are incorrect?"
+ Below we discuss wo points for such possibilities., Below we discuss two points for such possibilities.
+ First possibility. is that) our simulation might be overproducing the low-nmass galaxies at 4., First possibility is that our simulation might be overproducing the low-mass galaxies at $z>4$.
+ Recent observational and icoretical studies show that the low- hieh-: galaxies could. have lower SE elliciencies than the normal local galaxies (Wolfe&Chen2006:CinedinIxravtsov2010) owing to lower molecular hydrogen fractions in low-metallicity environments.," Recent observational and theoretical studies show that the low-mass, $z$ galaxies could have lower SF efficiencies than the normal local galaxies \citep{Wolfe.Chen:06,Gnedin.Kravtsov:10} owing to lower molecular hydrogen fractions in low-metallicity environments."
+" Lf this is true. our curren simulations might be overpredicting the p, at z29d."," If this is true, our current simulations might be overpredicting the $\rhostar$ at $z>4$."
+ However if we revise our SE model to account for this elfect. the SERD will also decrease together with pi. and we migh underpredict SERD instead.," However if we revise our SF model to account for this effect, the SFRD will also decrease together with $\rhostar$, and we might underpredict SFRD instead."
+ In the future we will revise our SE model to consider the LH» fraction in high-z galaxies. anc evaluate how strong this cllect would boe.," In the future we will revise our SF model to consider the $_2$ fraction in $z$ galaxies, and evaluate how strong this effect would be."
+ The second. possibility is that the IAL at high-z coul be dillerent from the local one., The second possibility is that the IMF at $z$ could be different from the local one.
+ All data in Figure 2. assume the Salpeter(1955). Το., All data in Figure \ref{fig:sf} assume the \citet{Salpeter:55} IMF.
+ As we emphasized in the [as paragraph of Section 2.. the amount of gas converted. into stars is a direct output of our simulations. therefore there are no uncertainties as to stellar IME except the instantaneous eas-recycling fraction. ic. the ratio between the gas expelled [rom stars (via stellar winds and supernovae) and the total initial stellar mass.," As we emphasized in the last paragraph of Section \ref{sec:method}, the amount of gas converted into stars is a direct output of our simulations, therefore there are no uncertainties as to stellar IMF except the instantaneous gas-recycling fraction, i.e., the ratio between the gas expelled from stars (via stellar winds and supernovae) and the total initial stellar mass."
+ In our simulations. we assume the Salpeter LME whose instantaneous gas-recycling fraction (:7) is 0.1 (Springel&Llerneuist 2003a)..," In our simulations, we assume the Salpeter IMF whose instantaneous gas-recycling fraction $\beta$ ) is $\sim$ 0.1 \citep{Springel.Hernquist:03}. ."
+ Lowe use the Chabrier(2003) IME that has a lower number of low-mass stars. ? increases to 70.2. which would reduce the simulated. p; by about1054.," If we use the \citet{Chabrier:03} IMF that has a lower number of low-mass stars, $\beta$ increases to $\sim$ 0.2, which would reduce the simulated $\rhostar$ by about."
+.. Note that 33 here only considers the σας-reeveling from. massive stars., Note that $\beta$ here only considers the gas-recycling from massive stars.
+ Lowe take into account of the contribution from long-lived. low-mass stars. the value of j will increase by a factor of a few.," If we take into account of the contribution from long-lived, low-mass stars, the value of $\beta$ will increase by a factor of a few."
+ However. this paper focuses on the high-z galaxies. and we can safely ignore the eas-reeveling from low-mass stars.," However, this paper focuses on the $z$ galaxies, and we can safely ignore the gas-recycling from low-mass stars."
+ Pherefore the change from Salpeter to Chabrier LME cannot fully account for the inconsistency between SERD and py., Therefore the change from Salpeter to Chabrier IMF cannot fully account for the inconsistency between SFRD and $\rhostar$.
+ The top-heavy Ιλ in. hieh-: galaxies has been speculated by many authors2010).. ancl it has à much higher value of 3.," The top-heavy IMF in $z$ galaxies has been speculated by many authors, and it has a much higher value of $\beta$ ."
+ Consequently. the resulting p; will decrease significantly: for a given SERD. and it may alleviate the discrepancy in. ρε.," Consequently, the resulting $\rhostar$ will decrease significantly for a given SFRD, and it may alleviate the discrepancy in $\rhostar$."
+ However. the current theoretical models and observational support for top-heavy EME are not very robust: for example. vanDokkum&Conroy(2010) suggested. abundant low-mass stars in high-z star-lorming galaxies. which argues against a top-heavy LAI.," However, the current theoretical models and observational support for top-heavy IMF are not very robust; for example, \citet{vanDokkum.Conroy:10} suggested abundant low-mass stars in $z$ star-forming galaxies, which argues against a top-heavy IMF."
+" lo summary. we find a good agreement between our self-consistent cosmological simulations and the observational estimates of SERD anc o, if we limit the comparison to galaxies with Ad,10h.TALS."," In summary, we find a good agreement between our self-consistent cosmological simulations and the observational estimates of SFRD and $\rhostar$ if we limit the comparison to galaxies with $\Mstar > 10^9\himsun$."
+ In. particular. the consisteney between our simulations and the observational estimates by Ixistlerctal.(2009)... Ouchictal.(2009) and Gonzálezctal.(2010) at 2>4 is very encouraging.," In particular, the consistency between our simulations and the observational estimates by \citet{Kistler.etal:09}, \citet{Ouchi.etal:09} and \citet{Gonzalez.etal:10} at $z>4$ is very encouraging."
+" Our simulations predict the existence of numerous low-mass galaxies. with. AL,«10°41AZ. and these low-mass galaxies. are not included in the current observational estimates of p.."," Our simulations predict the existence of numerous low-mass galaxies with $\Mstar < 10^9\himsun$, and these low-mass galaxies are not included in the current observational estimates of $\rhostar$."
+ This paper demonstrates that the current observational estimates and understanding of the formation of high-z galaxies are still uncertain., This paper demonstrates that the current observational estimates and understanding of the formation of $z$ galaxies are still uncertain.
+ To resolve this problem. we need future observations with increasing sensitivity aud reduced. uncertainties (e.g. JWST) to provide more robust constraints for the abundance of low-mass galaxies at 2>4. as well as the improvement in the theoretical modeling of galaxy formation. particularly on star formation and its feedback.," To resolve this problem, we need future observations with increasing sensitivity and reduced uncertainties (e.g., JWST) to provide more robust constraints for the abundance of low-mass galaxies at $z>4$, as well as the improvement in the theoretical modeling of galaxy formation, particularly on star formation and its feedback."
+ This work is supported in. part by the NSE grant AS'T-OS807491. NASA erant HST-AI-12143-01-A. National Aeronautics and Space Xediministration. under Grant/Cooperative Agreement No.," This work is supported in part by the NSF grant AST-0807491, NASA grant HST-AR-12143-01-A, National Aeronautics and Space Administration under Grant/Cooperative Agreement No."
+ NNNOSAISS7AA. issued by the Nevada NASA EPSCoR. program. and the President's Infrastructure Award from UNLV.," NNX08AE57A issued by the Nevada NASA EPSCoR program, and the President's Infrastructure Award from UNLV."
+ JIC acknowledges thesupport from University of Wentueky., JHC acknowledges thesupport from University of Kentucky.
+ This. research isalso supported bv the NSE through the TeraCirid. resources provided by the Texas Advanced Computing Center., This research isalso supported by the NSF through the TeraGrid resources provided by the Texas Advanced Computing Center.
+ Some numerical simulations and analysis have also been performed on the UNLY Cosmology. Cluster., Some numerical simulations and analysis have also been performed on the UNLV Cosmology Cluster.
+Carbon exists in several phases in the interstellar medium: dust. gaseous atoms/ions and gaseous molecules.,"Carbon exists in several phases in the interstellar medium: dust, gaseous atoms/ions and gaseous molecules."
+ The phase distribution of carbon is a fundamental characteristic of interstellar clouds that governs their physics., The phase distribution of carbon is a fundamental characteristic of interstellar clouds that governs their physics.
+ In the gaseous form carbon plays a crucial role in interstellar molecular chemistry ancl interstellar cloud cooling through fine structure lines., In the gaseous form carbon plays a crucial role in interstellar molecular chemistry and interstellar cloud cooling through fine structure lines.
+ Carbon that exists in the dust phase. as evidenced by spectral features. attributed {ο amorphous carbons (Dulev&Williams1981).. graphite (Draine1959) ancl polvevelic aromatic hvdrocarbons (Joblinetal.1992).. may be the dominant heat source [or the diffuse ISAT (Bakes&Tielens1994) and a large contributor to interstellar extinction (e.g.. ihe 2175 feature; Savage 1975: Draine 1989: Légger οἱ al.," Carbon that exists in the dust phase, as evidenced by spectral features attributed to amorphous carbons \citep{dw81}, graphite \citep{dra89} and polycyclic aromatic hydrocarbons \citep{jlm92}, may be the dominant heat source for the diffuse ISM \citep{bt94} and a large contributor to interstellar extinction (e.g., the 2175 feature; Savage 1975; Draine 1989; Légger et al."
+ 1989)., 1989).
+ In order to wnelerstand the physical processes and exünct(on properties of the ISM. we must know the distribution of carbon among ils possible phases.," In order to understand the physical processes and extinction properties of the ISM, we must know the distribution of carbon among its possible phases."
+ The dominant carbon ion in neutral interstellar clouds is €11. Unfortunately the abundance of this ion is extremely difficult to measure because it produces only very strong absorption features al 1036 and 1334 space((log fA = 2.088 and 2.234. respectively: Morton 2003) ancl an extremely weakinlersvsteni feature at 2325.4029 (log fA = -3.954: Morton 2003).," The dominant carbon ion in neutral interstellar clouds is C. Unfortunately the abundance of this ion is extremely difficult to measure because it produces only very strong absorption features at 1036 and 1334 (log $f\lambda$ = 2.088 and 2.234, respectively; Morton 2003) and an extremely weakintersystem feature at 2325.4029 (log $f\lambda$ = -3.954; Morton 2003)."
+ Complex stellar continua. multiple absorbing components and saturation effects severely limit the accuracy. of column density determinations from strong lines (errors can easily be a [actor of 2: Fitzpatrick 1997). so direct measurement of the weak intersvstem line is the best wav to find reliable C abundances in the neutral ISM.," Complex stellar continua, multiple absorbing components and saturation effects severely limit the accuracy of column density determinations from strong lines (errors can easily be a factor of 2; Fitzpatrick 1997), so direct measurement of the weak intersystem line is the best way to find reliable C abundances in the neutral ISM."
+ Because of the weakness of the intersvstem line. only 8 absorption measurements of the interstellar Co gas-phase abundance have been reported in the literature.," Because of the weakness of the intersystem line, only 8 absorption measurements of the interstellar C gas-phase abundance have been reported in the literature."
+ Seven of these abundance determinations were for diffuse sightlines (sightlines with total Ay<1.0: Hobbs. York Ogerle 1982: Cardelli et al.," Seven of these abundance determinations were for diffuse sightlines (sightlines with total $A_{V} \lesssim 1.0$; Hobbs, York Ogerle 1982; Cardelli et al."
+ 1993. 1996: Sofia et al.," 1993, 1996; Sofia et al."
+" 1997): the eighth was through a translucent sightline toward ID 24534 CA,= 1.44: Solia. Fitzpatrick Meyer 1998)."," 1997); the eighth was through a translucent sightline toward HD 24534 $A_{V} = 1.44$ ; Sofia, Fitzpatrick Meyer 1998)."
+ If the, If the
+~ 2 20062~ and barvouic processes driviug this rapid evolution is therefore central to our nuderstanuding of ealaxy formation aud to informing cosmological sinilations.,$\sim$ $-$ $z \sim$ $z \sim$ and baryonic processes driving this rapid evolution is therefore central to our understanding of galaxy formation and to informing cosmological simulations.
+ At the relevant epochs. kev spectral diagnostic features are redshüfted iuto the near-intrared.," At the relevant epochs, key spectral diagnostic features are redshifted into the near-infrared."
+ In receut veurss a ΠΙΟ of surveys have therefore begun to systematically probe ligh-+vedshift popuations with near-infrared spectroscopy2008).," In recent years, a number of surveys have therefore begun to systematically probe high-redshift populations with near-infrared spectroscopy."
+. Our recently completed SINS (S)ectroscopie Thuaging iu the Near-infrared with SINFONT) survey las combined the resolving power of S-LOm class telescoTES with high-resolution iutegral field spectrograph to stidy Thie detailed internal processes at work witjin massive. 4AWemar-formine galaxies atal.1092: see also related workby ).," Our recently completed SINS (Spectroscopic Imaging in the Near-infrared with SINFONI) survey has combined the resolving power of 8-10m class telescopes with high-resolution integral field spectrographs to study the detailed internal processes at work within massive, star-forming galaxies at; see also related work by )."
+ Iu this paper. we combine the spectra of eealaxies detected in eenission to study the average spectral propvties of forming galaxies at rofData)).," In this paper, we combine the spectra of galaxies detected in emission to study the average spectral properties of star-forming galaxies at \\ref{Data}) )."
+ Iu particular. we report the discoviery of broad enission lines in these galaxies refResults)} aud interpret this high-velocity wari gas as arising either in large-scale galactic winels driven by the high star formation rates (SFR) in these galaxies or in the broad-line regions (BLR) surrounding active ealactic nuclei (ACN).," In particular, we report the discovery of broad emission lines in these galaxies \\ref{Results}) ) and interpret this high-velocity warm gas as arising either in large-scale galactic winds driven by the high star formation rates (SFR) in these galaxies or in the broad-line regions (BLR) surrounding active galactic nuclei (AGN)."
+ We explore the iu]ylications of both scenarios in and conclude in ouch..., We explore the implications of both scenarios in \\ref{Discussion} and conclude in \\ref{Conclu}. .
+program (Cherenkov photons m 300150 nin).,program (Cherenkov photons in 300–450 nm).
+ In the Figure ?? we show time distributions of and angular aud production Leight distributions of plotous correspoucding to photoelectrous for the αι x 3 ii detector at (501. 0 11) pointing vertically with angular acceptance within 20 nirad (related to the Figure 3)).," In the Figure \ref{fig:pe:0:0} we show time distributions of photo--electrons and angular and production height distributions of photons corresponding to photo–electrons for the 3 m x 3 m detector at (–50 m, 0 m) pointing vertically with angular acceptance within 20 mrad (related to the Figure \ref{fig:0:0}) )."
+ The difference due to transition from the ‘emitted photon case to the case can be seen by comparison of ordinate axes between Figures 9 aud 2? for c) aud d) Tu the Figure ?? we preseut the photoelectron distributions for detector poiuting to (LO πας. 10 iirad) direction to the EAS direction and including the umon signal (related to the Figure ??)).," The difference due to transition from the emitted photon' case to the photo--electron' case can be seen by comparison of ordinate axes between Figures \ref{fig:0:0} and \ref{fig:pe:0:0}
+ for c) and d) In the Figure \ref{fig:pe:10:10} we present the photo–electron distributions for detector pointing to (10 mrad, 10 mrad) direction to the EAS direction and including the muon signal (related to the Figure \ref{fig:10:10}) )."
+ For the discussed. problem of Cherenkov helt from fast nimous the 2 us gap secu in the Figure ??7cc is very The Fieures ??cc aud οτου (Af distributions for photoelectrons) differ quite significauth., For the discussed problem of Cherenkov light from fast muons the 2 ns gap seen in the Figure \ref{fig:pe:10:10}c c is very The Figures \ref{fig:pe:0:0}c c and \ref{fig:pe:10:10}c c $\Delta t$ distributions for photo–electrons) differ quite significantly.
+ However. phlotoniultiplier anode sigual does not reproduce these patterns accurately.," However, photomultiplier anode signal does not reproduce these patterns accurately."
+ Hore we present results of sienal convolution for two photomultiplers: PITILIPS. XP2020 and Waimamatsu 12083. (, Here we present results of signal convolution for two photomultipliers: PHILIPS XP2020 and Hamamatsu H2083. (
+The XP2020 data afterhilipsHand aud Tamamatsu 12083 after icra).,The XP2020 data after and Hamamatsu H2083 after ).
+ We use following parametrization of the anode inipulse corresponding to one photoelectron where R5(f) is the time dependent (uegative) anode signal corresponding to ó impulse one G is the actual gain (current amplification). f£ is time (see below for more details). i. ση are paranieters related to the pulse shape.," We use following parametrization of the anode impulse corresponding to one photo–electron : where $R_{\delta}(t)$ is the time dependent (negative) anode signal corresponding to $\delta$ impulse – one photo--electron, $G$ is the actual gain (current amplification), $t$ is time (see below for more details), $m$, $\sigma_{\rm R}$ are parameters related to the pulse shape."
+ The 5(f) has following properties: The time between the photoclectron emission from the cathode auc time of the maxim of sienal fae is called the transit time auc its average value depends on photomultiplier. high voltage and divider.," The $R_{\delta}(t)$ has following properties: The time between the photo–electron emission from the cathode and time of the maximum of signal $t_{\qqq{max}}$ is called the transit time and its average value depends on photomultiplier, high voltage and divider."
+ In this calculation we sot the average transit the to be 30 ns for both photommltiplicrs (although it is different in the real case. e.g. about 39 us for NP2020).," In this calculation we set the average transit time to be 30 ns for both photomultipliers (although it is different in the real case, e.g. about 39 ns for XP2020)."
+ Jitter is the variability of that time., Jitter is the variability of that time.
+ Jitter has the Gaussian distribution with σεjitter = 0.25 us., Jitter has the Gaussian distribution with $\sigma_{\qqq{t\_ jitter}}$ = 0.25 ns.
+" We asstune another fluctuation of transit time due to the differeut photoelectron emission place iu the cathode to be of Gaussian shape with of44,4; = 0.25 where rise time is the time between aud of maxima value. 1 aud oR are parameters of equation 1.."," We assume another fluctuation of transit time due to the different photo–electron emission place in the cathode to be of Gaussian shape with $\sigma_{\qqq{t\_ cat}}$ = 0.25 where rise time is the time between and of maximum value, $m$ and $\sigma_{\rm R}$ are parameters of equation \ref{eq:PMsignal}."
+ Iu the Figure ὃ we preseut the simulated shapes of anode signals (reversed to positivo) for photoniultipliers PITILIPS XP2020 and Wamamatsu II2085 for input photoelectron distributions presented in the Figures ??cc aud ??cc. Units are arbitrary. since. effectively. this is further electronics depeudent.," In the Figure \ref{fig:anode} we present the simulated shapes of anode signals (reversed to positive) for photomultipliers PHILIPS XP2020 and Hamamatsu H2083 for input photo–electron distributions presented in the Figures \ref{fig:pe:0:0}c c and \ref{fig:pe:10:10}c c. Units are arbitrary, since, effectively, this is further electronics dependent."
+ One can notice that at low discrimination level the signalsC» with nmon aud EM Chereukov light[m] last longeroO than signalsoO without muon., One can notice that at low discrimination level the signals with muon and E–M Cherenkov light last longer than signals without muon.
+ Iu both cases (1.0. NP2020 and Tamaimatsu 112083) the signaloO for the case with the fast umon iu the field of view starts 2 us earlier than the signal form EAL Cherenkov This fact (time of start of the signal iu one detector) does not discriminate between the no.mmon and muon cases., In both cases (i.e. XP2020 and Hamamatsu H2083) the signal for the case with the fast muon in the field of view starts 2 ns earlier than the signal form E–M Cherenkov This fact (time of start of the signal in one detector) does not discriminate between the no–muon and muon cases.
+ To see the effect of fast nion(examining the start time) if is necessary to examine the start time in a number of detectors for the same event., To see the effect of fast muon(examining the start time) it is necessary to examine the start time in a number of detectors for the same event.
+ Du most cases (also experimental) it is, In most cases (also experimental) it is
+Among the ILCGs. we have selected the groups with redshifts 2<0.045. having a surface brightness jt;<24.4.,"Among the HCGs, we have selected the groups with redshifts $z\leq0.045$, having a surface brightness $\mu_G\leq24.4$."
+ These criteria provided us with a statistically complete sample of 283 ealaxies in 65 eroups., These criteria provided us with a statistically complete sample of 283 galaxies in 65 groups.
+ We have obtaimed medium resolution spectroscopy for 238 of these ealaxies., We have obtained medium resolution spectroscopy for 238 of these galaxies.
+ The spectra of Tl galaxies come from previous observations made by 2000.2004).," The spectra of 71 galaxies come from previous observations made by \citet{coz98,coz00,coz04}."
+. The remaining 167 galaxies were observed by our group using four cilferent telescopes: the 2.5m in “El Roque de los Muchachos? (RM. Spain). the 2.2m in Calar Alto(CAILA?.. Spain). the 2.12m in San Pedro Márrtir (SPM. Mexico) and the 1.5m telescope in Sierra Nevada Observatory (OSN. Spain).," The remaining 167 galaxies were observed by our group using four different telescopes: the 2.5m in “El Roque de los Muchachos” (RM, Spain), the 2.2m in Calar Alto, Spain), the 2.12m in San Pedro Márrtir (SPM, Mexico) and the 1.5m telescope in Sierra Nevada Observatory (OSN, Spain)."
+ For all the galaxies. broad components search and activity classification were done only aller template subtraction. (o correct [or absorption features produced by (he underlviug stellar population.," For all the galaxies, broad components search and activity classification were done only after template subtraction, to correct for absorption features produced by the underlying stellar population."
+ Detailed of the template subtraction method used can be found in Coziol 2004).," Detailed of the template subtraction method used can be found in \citet{coz98,coz04}."
+. Preliminary results have already been published in Martínezetal.(2007)., Preliminary results have already been published in \citet{mar07}.
+. Complete description of observations. reduction and analvsis method will be published elsewhere.," Complete description of observations, reduction and analysis method will be published elsewhere."
+ For (he UZC-CG sample. we have collected spectra from three spectroscopic archives: the Sloan Digital Sky Survey (5DSS-DRA). the Z-Machine and the FAST Spectrograph Archives.," For the UZC-CG sample, we have collected spectra from three spectroscopic archives: the Sloan Digital Sky Survey (SDSS-DR4), the Z-Machine and the FAST Spectrograph Archives."
+ We have [ound spectra for all the galaxy. menmbers of 215 groups (720 galaxies): 210 spectra are [rom the SDSS. 187 from FAST and 323 [rom Z-Machine (Martinezetal.2006).," We have found spectra for all the galaxy members of 215 groups (720 galaxies): 210 spectra are from the SDSS, 187 from FAST and 323 from Z-Machine \citep{mar06}."
+. Becatse the Z-Machine spectra have too low S/N ratios to measure broad components. we restrict our analvsis to the 397 spectra found in the SDSS and FAST databases.," Because the Z-Machine spectra have too low S/N ratios to measure broad components, we restrict our analysis to the 397 spectra found in the SDSS and FAST databases."
+ Spectra from the SDSS survev were template subtracted using Laoetal.(2005.1105) eigenspectra and (heir PCA method., Spectra from the SDSS survey were template subtracted using \citet[H05]{hao05} eigenspectra and their PCA method.
+ No correction was applied to the galaxies in the FAST sample. due to the non availability of suitable spectra to be used as templates.," No correction was applied to the galaxies in the FAST sample, due to the non availability of suitable spectra to be used as templates."
+ Emission lines were found in 153 of the 238 galaxies in the HCG sample )) and 274 of the 397 galaxies )) in the UZC-CG sample., Emission lines were found in 153 of the 238 galaxies in the HCG sample ) and 274 of the 397 galaxies ) in the UZC-CG sample.
+ The identilication of broad components was done by fitting à multi-components Gaussian on (he emission lines. using the ILAF task NGAUSSFIT.," The identification of broad components was done by fitting a multi-components Gaussian on the emission lines, using the IRAF task NGAUSSFIT."
+" The FWIIM of [SII] (or (OUT) when the [SII] lines were too faint or noisy) have been used to model the narrow components of the IL, and [NH] lines.", The FWHM of [SII] (or [OIII] when the [SII] lines were too faint or noisy) have been used to model the narrow components of the $_{\alpha}$ and [NII] lines.
+ When a broadfourth component was necessary. it was centered on IL.," When a broadfourth component was necessary, it was centered on $_{\alpha}$ ."
+ A 4? criterion. as described in I105.," A $\chi^2$ criterion, as described in H05,"
+Seurches for galaxy overdensities near these objects have resulted in detections of large-scale structures and protoclusters up to 5.3 (eg.2222200)..,"Searches for galaxy overdensities near these objects have resulted in detections of large-scale structures and protoclusters up to $z=5.2$ \citep[e.g.,][]{LeFevre1996,Pentericci2000,Kurk2000,Kurk2004,Best2003,Overzier2006,Venemans2007,Overzier2008}."
+ We have performed an infrared survey of the fields of 8 powerful radio galaxies with redshifts in the range 1.7xz2.6 using the wide-field near-infrared imager HAWK-I on the VLT., We have performed an infrared survey of the fields of 8 powerful radio galaxies with redshifts in the range $1.7\le z \le 2.6$ using the wide-field near-infrared imager HAWK-I on the VLT.
+ The aim of the survey is to identify protoclusters and study the member galaxies and their evolution over this redshift range., The aim of the survey is to identify protoclusters and study the member galaxies and their evolution over this redshift range.
+ The HAWK-I data on the two lowest redshift targets 220andMRCOIS6—252 are presented in ?.., The HAWK-I data on the two lowest redshift targets $1017-220$ and $0156-252$ are presented in \citet{Galametz2010}.
+ In this paper we present HAWK-I data for the 6 highest redshift targets. ranging from 2.2«zx2.6.," In this paper we present HAWK-I data for the 6 highest redshift targets, ranging from $2.2 \le z \le 2.6$."
+ In Section 2. we describe the observations and data reduction., In Section \ref{obs_dr} we describe the observations and data reduction.
+ Sections 3 — 5 lay out the evidence that 3 out of 6 of the radio galaxies He within galaxy protoclusters., Sections \ref{identify} – \ref{stat} lay out the evidence that 3 out of 6 of the radio galaxies lie within galaxy protoclusters.